Spectroscopy and Chemistry of Interstellar Ice Analogues

Spectroscopy and Chemistry of Interstellar Ice Analogues
Spectroscopy and Chemistry of
Interstellar Ice Analogues
Spectroscopy and Chemistry of Interstellar Ice Analogues – Jordy Bouwman
Thesis Universiteit Leiden - Illustrated - With summary in Dutch - With references
ISBN/EAN 978-90-9025686-3
Printed by Ipskamp Drukkers
Cover by Ruud Engelsdorp
This work is part of the research programme of the Foundation for Fundamental Research on Matter (FOM), which is part of the Netherlands
Organisation for Scientific Research (NWO).
Spectroscopy and Chemistry of
Interstellar Ice Analogues
PROEFSCHRIFT
ter verkrijging van
de graad van Doctor aan de Universiteit Leiden,
op gezag van de Rector Magnificus prof. mr. P. F. van der Heijden,
volgens besluit van het College voor Promoties
te verdedigen op dinsdag 12 oktober 2010
klokke 13.45 uur
door
Jordy Bouwman
geboren te Haarlem
in 1979
Promotiecommissie:
Promotor:
Copromotor:
Prof. dr. H. V. J. Linnartz
Dr. L. J. Allamandola
Overige Leden:
Prof. dr. K. Kuijken
Prof. dr. A. G. G. M. Tielens
Prof. dr. M. R. S. McCoustra
Prof. dr. J. Oomens
Dr. H. M. Cuppen
NASA Ames Research Center
Heriot-Watt University
FOM Rijnhuizen
Contents
1 Introduction
1.1 Astrochemistry . . . . . . . . . . . . . . . . . . . . . . . .
1.2 The interstellar cycle of matter . . . . . . . . . . . . . . . .
1.3 Mid-IR absorption bands – Interstellar ices . . . . . . . . . .
1.3.1 Composition of interstellar ices . . . . . . . . . . .
1.3.2 Ice formation and grain chemistry . . . . . . . . . .
1.4 Mid-IR emission bands – Polycyclic Aromatic Hydrocarbons
1.4.1 The PAH building block – Carbon . . . . . . . . . .
1.4.2 The origin of interstellar PAHs . . . . . . . . . . . .
1.4.3 PAHs in interstellar ices? . . . . . . . . . . . . . . .
1.5 Laboratory spectroscopic ice studies . . . . . . . . . . . . .
1.5.1 Mid-IR ice spectroscopy . . . . . . . . . . . . . . .
1.5.2 Near-UV/VIS absorption ice spectroscopy . . . . . .
1.6 Outline of this thesis . . . . . . . . . . . . . . . . . . . . .
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1
1
3
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I Mid-IR absorption spectroscopy
19
2 Band profiles and band strengths in mixed H2 O:CO ices
2.1 Introduction . . . . . . . . . . . . . . . . . . . . . .
2.2 Experiment and data analysis . . . . . . . . . . . . .
2.3 Results . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Influence of CO on water bands . . . . . . .
2.3.2 Influence on the CO band . . . . . . . . . .
2.4 Discussion . . . . . . . . . . . . . . . . . . . . . . .
2.5 Conclusions . . . . . . . . . . . . . . . . . . . . . .
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V
Contents
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
3.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . .
3.2 Astronomical observations and analysis . . . . . . . . .
3.2.1 Local continuum . . . . . . . . . . . . . . . . .
3.2.2 Template . . . . . . . . . . . . . . . . . . . . .
3.2.3 NH3 ice column densities and abundances . . .
3.3 Laboratory work and analysis . . . . . . . . . . . . . . .
3.4 Comparison between astronomical and laboratory data .
3.4.1 8–10 µm range . . . . . . . . . . . . . . . . . .
3.4.2 The 3 and 6 µm ranges . . . . . . . . . . . . . .
3.4.3 Nitrogen ice inventory . . . . . . . . . . . . . .
3.5 Conclusion . . . . . . . . . . . . . . . . . . . . . . . .
3.6 Appendix . . . . . . . . . . . . . . . . . . . . . . . . .
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39
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4 IR spectroscopy of PAH containing ices
4.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2 Experimental technique . . . . . . . . . . . . . . . . . . . .
4.3 PAH:H2 O spectroscopy . . . . . . . . . . . . . . . . . . . .
4.4 PAH ice photochemistry . . . . . . . . . . . . . . . . . . .
4.4.1 PAH:H2 O photoproducts . . . . . . . . . . . . . . .
4.4.2 Concentration effects and time dependent chemistry
4.4.3 Ionization efficiency in CO ice . . . . . . . . . . . .
4.4.4 Temperature effects . . . . . . . . . . . . . . . . . .
4.5 The non-volatile residue . . . . . . . . . . . . . . . . . . .
4.6 Astrophysical implications . . . . . . . . . . . . . . . . . .
4.6.1 High-mass protostars . . . . . . . . . . . . . . . . .
4.6.2 Low-mass protostars . . . . . . . . . . . . . . . . .
4.6.3 PAH contributions to the 5–8 µm absorption . . . . .
4.7 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . .
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II Near-UV/VIS absorption spectroscopy
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
5.1 Introduction . . . . . . . . . . . . . . . . . . . . . .
5.2 Experimental . . . . . . . . . . . . . . . . . . . . .
5.3 Spectroscopic assignment . . . . . . . . . . . . . . .
5.4 Chemical evolution of the ice . . . . . . . . . . . . .
5.5 Astrophysical implications . . . . . . . . . . . . . .
5.6 Conclusion . . . . . . . . . . . . . . . . . . . . . .
VI
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109
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Contents
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
6.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . .
6.2 Experimental technique . . . . . . . . . . . . . . . . . . . .
6.3 Band assignments and band strength analysis . . . . . . . .
6.3.1 Neutral pyrene bands . . . . . . . . . . . . . . . . .
6.3.2 Pyrene cation bands . . . . . . . . . . . . . . . . .
6.3.3 HCO bands in Py:CO . . . . . . . . . . . . . . . . .
6.3.4 The 400 nm band carrier . . . . . . . . . . . . . . .
6.3.5 The 405 nm band carrier . . . . . . . . . . . . . . .
6.3.6 Broad absorption feature . . . . . . . . . . . . . . .
6.4 Py:H2 O ice photochemistry at different temperatures . . . .
6.5 Astrochemical Implications . . . . . . . . . . . . . . . . . .
6.6 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . .
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7 Ionization of PAHs in interstellar ices
7.1 Introduction . . . . . . . . . . . . . .
7.2 Experimental technique . . . . . . . .
7.3 PAH:H2 O spectroscopy . . . . . . . .
7.3.1 Anthracene (C14 H10 ) . . . . .
7.3.2 Pyrene (C16 H10 ) . . . . . . .
7.3.3 Benzo[ghi]perylene (C22 H12 ) .
7.3.4 Coronene (C24 H12 ) . . . . . .
7.4 PAH ionization rates . . . . . . . . .
7.5 Astrophysical implication . . . . . . .
7.6 Conclusions . . . . . . . . . . . . . .
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8 Future challenges
161
Bibliography
165
Nederlandse samenvatting
173
Publications
179
Curriculum vitae
181
Nawoord
183
VII
CHAPTER 1
Introduction
The pressures in space are much lower than one can reach in the best vacuum chamber in
a laboratory on Earth and the temperatures vary from extremely high to close to absolute
zero. Despite these extreme circumstances there is a surprisingly active chemistry which
enriches the vast regions in space, leaving a large puzzle for mankind to solve.
Most molecules detected in the interstellar medium (ISM) are unambiguously identified by their rovibrational (infrared), or purely rotational (microwave) fingerprint absorption or emission spectra. One particular family of molecules — the so-called Polycyclic
Aromatic Hydrocarbons (PAHs) — is detected as a class by its characteristic mid-infrared
(mid-IR) emission spectrum. Although these molecules have not been uniquely identified,
because of their common spectral signature, their presence in photon-dominated regions
(PDRs) is now widely accepted in the astrochemical community.
Besides gas phase species, molecules are also detected in solid form, as interstellar
ices. Ices are formed in cold and dark regions in space, known as molecular clouds, by
accretion of gas phase species on cold carbonaceous or silicate dust grains. The thin
layers of ice contain rather simple molecules, such as H2 O, CO, CO2 , CH3 OH, CH4 ,
and NH3 . The constituents of icy grain mantles are further energetically processed by
heat, cosmic rays, or ultraviolet (VUV) radiation, leading to more complex molecules.
Interstellar ices are now regarded as important catalytic sites for the formation of complex
(organic) molecules during the evolution of an interstellar cloud and are considered crucial
in astrochemistry.
This thesis describes laboratory and observational studies which are aimed to understand physical interactions and abundances of species in, and the VUV induced chemical
evolution of, interstellar ices in a laboratory setting using mid-IR and near-ultraviolet/ visible (near-UV/VIS) spectroscopic techniques. The remainder of the introduction is used
to put the thesis work into context.
1.1 Astrochemistry
The formation and detection of polyatomic molecules in the interstellar medium had long
been unexpected because of the harsh UV fields and low densities (∼ 1–102 molecules
1
1 Introduction
Diffuse clouds gather into
large molecular clouds
In the interstellar medium gas and
dust is exposed to shocks and the
interstellar radiation field, shattering
and gas-phase reactions alter the dust
and polycyclic aromatic hydrocarbons
Collapse and fragmentation of
the molecular cloud, resulting
in dense proto-stellar cores
At the end of their life massive
stars inject gas and dust into the
interstellar medium through
supernovea, introducing shocks
Old stars expelling their outer
layers, enriching the interstellar
medium with gas and dust, including polycyclic aromatic hydrocarbons formed in the ejecta
A long lived main sequency
star with a planetary system
Low and intermediate mass stars
form disks where planets can form
Figure 1.1 Cartoon of the Galactic life cycle. After Steven Simpson (Verschuur 1992, Sky
& Telescope Magazine), by Christiaan Boersma.
cm−3 ). However, in 1937 the first molecules, CH, CN, and CH+ were detected in diffuse
clouds [Swings & Rosenfeld 1937]. The detection of only transient species confirmed
the idea that the unfavorable conditions would preclude the presence of more complex
chemistry. The detection of more complex molecules in the ISM such as NH3 and H2 CO
in the 1960s opened up a new field in astronomy, astrochemistry, in which the abundances
and reactions of chemical elements and molecules, and their interaction with light are
studied. Up to now, as many as 152 molecules1 have been detected in the gas of interand circumstellar clouds and every year some new species are detected. Amongst the
detected molecules are simple species, such as H2 and CO, but also rather complex and
exotic species, such as HC11 N [Bell et al. 1997] with the largest unambiguously detected
molecules being C60 and C70 . The detection of a large variety of species in the strongly
UV processed medium implies that chemical reactions are very efficient. It is mostly ionmolecule reactions that are responsible for the high production rates of these molecular
species in the gas. In cold regions, such as dense clouds, chemistry is now known to
proceed via grain catalyzed reactions in which species released from interstellar ices play
a key role.
New ground-based and space-borne observatories with improved sensitivity and spectral resolution combined with advances in laboratory techniques shed new light on the
molecular diversity. The detected species continue giving us insight in the complex chemistry that takes place in the vast regions of space and perhaps even clues to the formation
of life on Earth, or even life outside of our own solar system.
1 http://www.astrochymist.org
2
1.2 The interstellar cycle of matter
1.2 The interstellar cycle of matter
Although what triggered the formation of the first stars in the Universe after the occurrence of the big bang about 14 billion years ago is still a big mystery, the life cycle of
low-mass stars, such as our own sun, is now quite well understood [e.g. Evans 1999, van
Dishoeck 2004, and references therein]. Disregarding the birth of the first stars, the evolution of gas and dust in the ISM from stellar birth to death can be depicted as a cyclic
event as seen in Fig. 1.1. The building blocks of the newly formed stars are the remnants of the old dead stars; the diffuse interstellar medium is enriched by its previous
inhabitants. Stellar remnants, however, are mostly destroyed by the omnipresent strong
ultraviolet (UV) radiation, leaving only heavy elements, large molecules such as PAHs,
and dust grains intact. New stars are formed from these basic ingredients and the complex
chemistry involved in star-formation starts all over again.
The process of low-mass star formation is schematically displayed in Fig. 1.2. In
the first stage, the diffuse medium is disturbed by a process, such as a stellar wind or
a supernova explosion. This causes dense clouds to form out of the material in the diffuse medium. Once formed, these dense cloud cores, mainly consisting of hydrogen,
helium, heavier elements, dust and some larger molecules, are held together by gravitational forces. The densities in these clouds reach a point at which the core of the
Figure 1.2 Schematic illustration of the different stages of low-mass star formation. Figure
taken from Visser [2009].
3
1 Introduction
cloud is completely shielded from intense UV radiation and molecules can form. The
temperatures in these dense clouds are low (T ∼10 K) and the densities rather high by
interstellar standards (∼ 104 –105 molecules cm−3 ), causing molecules to freeze out efficiently on nanometer sized dust particles or possibly on large PAH molecules or clusters
of PAH molecules. Thin layers of ice which act as catalytic sites for chemical reactions
are formed. Ice abundances, formation, and chemistry will be described in more detail in
§1.3.
Within the dense molecular clouds, cores of even higher densities (>105 molecules
−3
cm ) are formed. If the density in such a core gets high enough, the core will collapse
under its own gravity, forming a so-called protostellar core, i.e. a region of the cloud that
will eventually become a star. The collapse releases a large amount of energy and the pressure building up in the core prevents it from collapsing further. At this stage, molecules
play a key role in the process of star formation; they convert translational energy via collisions into IR radiation which is emitted at the molecule’s specific wavelengths. Some of
this radiation can escape the collapsing cloud, resulting in efficient cooling and a continuation of the collapse of the core.
In the next stage, the protostar starts losing angular momentum by expelling mass in
bipolar outflows. Additionally, a protoplanetary disk is formed around the central object
from which material continues to accrete onto the protostar. In this disk small grains
coagulate, forming larger and larger rocks and eventually planets. The outside of the disk
is processed by the strong UV irradiation from the new born star and becomes heated and
chemically processed. The center of the disk, however, remains cold and the chemical
evolution of matter in this part of the disk will be dominated by ice grain chemistry. For
low-mass stars, the disk will slowly evolve into a planetary system such as our own.
The formation process of high-mass stars is not yet fully understood, but most likely
has many similarities to the formation of low-mass stars. The final stages of the lifecycle
of both high- and low-mass stars, on the other hand, are well understood. At the end of its
life, the star enriches the interstellar medium by expelling its contents into its surroundings. Stars with a mass smaller than 2.5 Solar masses (M ≤ 2.5M⊙ ) such as our own
Sun expel their mass in relatively gentle stellar winds, the so-called protoplanetary nebula
phase, after which only the hot core of the star will remain; a white dwarf. Stars with a
large mass (M ≥ 2.5M⊙ ) end their life in a less gentle manner. They return their mass to
the ISM in a violent event, a so-called supernova explosion, which can again trigger the
formation of new stars as described above.
1.3 Mid-IR absorption bands – Interstellar ices
The presence of ices in the interstellar medium was already proposed in 1937, even before
the detection of the first interstellar molecule [Eddington 1937]. The detection of an
interstellar ice absorption feature was a fact in 1973, nearly four decades after the presence
of ices was proposed. A strong and broad mid-IR absorption band located at ∼3 µm
was detected and assigned to the H2 O ice stretching mode [Gillett et al. 1973]. The
spectroscopic signatures of interstellar ices fall in the mid-IR as absorption profiles which
4
1.3 Mid-IR absorption bands – Interstellar ices
are superimposed on the black body radiation curve of a background star, or embedded
object. Since the molecules are confined within the ice, they do not have translational, nor
rotational degrees of freedom and absorption of a mid-IR photon by a molecule results in
vibration of the molecule only. While interstellar species in the gas phase can be detected
in absorption or emission, ices are only detected in absorption. This comes from the fact
that the temperature corresponding to mid-IR radiation is higher than the temperature of
the ice.
Interstellar ices have been detected either using ground-based, airborne, or space
based observatories. Ground based mid-IR observatories, such as the powerful Very Large
Telescope (VLT), have a limited spectral window because nearly half of the mid-IR spectrum is blocked by telluric absorptions, primarily H2 O and CO2 . Sophisticated observatories were built to extend wavelength coverage, push the detection limit and to obtain
higher resolution spectra. First, the Kuiper Airborne Observatory (KAO) was constructed.
Observations with the KAO were conducted at high altitude (40,000 to 45,000 feet), well
above most of the H2 O in the atmosphere. This opened up a very important 5 to 10 µm
portion of the infrared fingerprint region [Haas et al. 1995]. The combination of airborne
with ground based observations provided the first access to nearly all of the mid-IR spectrum for a handful of objects. By the early 1990s, interstellar ices were known to be
water-rich mixtures containing species such as CH3 OH, NH3 , H2 CO, etc. The complete
mid-IR spectrum of the cosmos was opened up with the launch of the Infrared Space Observatory (ISO), an observatory that revolutionized our understanding of interstellar ices.
Free of telluric absorptions, the eyes of ISO revealed many secrets of ices in dense clouds
and around star forming regions. The number of detected interstellar ices nearly doubled.
While ISO probed quiescent lines of site as well as star forming regions, due to its low
sensitivity, however, ISO was limited to observing bright, high-mass, star-forming regions
[e.g., van Dishoeck 2004, and references therein]. Its successor, NASA’s very sensitive
Spitzer Space Telescope, opened up the window of opportunity further. It offered the high
sensitivity needed to observe faint objects such as low-mass protostars, without being limited by the transmission of the Earth’s atmosphere [Chapter 3, Pontoppidan et al. 2008,
Boogert et al. 2008, Öberg et al. 2008]. These very successful observatories offered a
sensitive view into the kitchen of newborn high- and low-mass stars.
1.3.1 Composition of interstellar ices
It is now established that water is the first molecule to form and freeze out on interstellar
grains in the evolution from a diffuse cloud to a dense cloud and that H2 O is the most
abundant species in ice toward most sources [e.g. Sonnentrucker et al. 2008]. Typical
Spitzer absorption spectra combined with L and M band VLT data toward two low-mass
protostars with the identified ice absorption bands marked out is shown in Fig. 1.3. The
ice absorption profiles are always accompanied by a feature at 10 µm which is typical
for the silicate Si–O stretching mode originating from the grain core. Table 1.1 gives
an overview of ice abundances with respect to H2 O ice detected towards the high-mass
object W33A.
5
1 Introduction
Besides the identified species indicated in Fig. 1.3, the detections of some other
molecules based on absorptions in the 5 to 8 µm spectral region have been suggested and
are given in parenthesis in Table 1.1. The dominant absorption profile in this region is
the 6 µm H2 O bending mode on which a substructure is superimposed. Besides the H2 O
bending mode and an absorption band at 7.68 µm, which can confidently be attributed
to the CH4 deformation mode. Assignments of other bands in this spectral region remain controversial. For example, the detection of species such as formaldehyde (H2 CO),
formic acid (HCOOH) and the ammonium ion (NH+4 ) have been claimed. Experiments
on the formation route of these molecules indeed point out that these molecules are the
likely formed under interstellar conditions and thus these species are plausible carriers of
these absorption bands [Fuchs et al. 2009, Ioppolo et al. 2010]. An absorption located
at 6.2 µm has been tentatively assigned to the CC stretching mode of aromatic molecules
trapped in the interstellar ice based on proximity to an interstellar emission band attributed
to aromatic species [Keane et al. 2001a]. Experimental data on the spectroscopy of these
species in interstellar ice analogues, however, is lacking in the literature. Chapter 4 deals
with the spectroscopy of aromatic molecules in ices and their possible contribution to
several absorption features in the 5 to 9 µm region.
Figure 1.3 Spitzer infrared absorption spectrum combined with L and M band observations of low-mass embedded protostars B5 IRS1 (top, multiplied by factor of 5 for clarification) and HH46 IRS (bottom). Identifications and possible identifications are indicated.
Spectrum is adopted from Boogert et al. [2004]
6
1.3 Mid-IR absorption bands – Interstellar ices
Table 1.1 Ice abundances with respect to H2 O ice towards the high-mass protostar W33A
taken from Gibb et al. [2000]. The NH+4 abundance is taken from Boogert & Ehrenfreund
[2004].
Species
H2 O
CO (polar)
CO (non-polar)
CO2 (polar)
CO2 (non-polar)
CH4
CH3 OH
Abundance
% of H2 O
100
6
2
11
2
1.5
18
Species
(HCOOH)
(H2 CO)
(NH3 )
(NH+4 )
OCN−
(SO2 )
OCS
Abundance
% of H2 O
7
6
15
12
3.5
2.4
0.2
The detection of NH3 ice has been claimed in some studies [e.g. Gibb et al. 2000, Lacy
et al. 1998] and upper limits of its abundance towards massive YSO’s have been reported
in others [e.g. Dartois & d’Hendecourt 2001]. Detections towards low-mass Young Stellar
Objects (YSO’s), however, remain controversial [Taban et al. 2003]. Most of the NH3
vibrations overlap with other prominent bands in the spectrum. The most isolated band,
the umbrella mode at ∼9 µm, overlaps with the strong 10 µm silicate absorption band.
The detection of NH3 in low-mass star forming regions is confirmed and investigated in
detail in Chapter 3 of this thesis.
Apart from the identification of frozen out species, mid-IR absorption spectra also
allow one to obtain information on physical properties of the ice, such as ice temperature,
degree of mixing, and interactions between species. Precise peak positions and band profiles directly reflect the composition and complex physical interplay between the species
in ices. This allows observers to discriminate, for example, between polar ices (H2 O-rich)
and non-polar ices (H2 O-poor) [Sandford et al. 1988] ice composition, in turn, reflects the
formation mechanisms and accretion history of molecules on cold grains. To this end, the
interaction between CO and H2 O in binary ices and the effect of mixing rations on band
shapes and band strengths is studied in detail in Chapter 2. A similar extensive laboratory
study on the effects of mixing H2 O, NH3 , CH3 OH , CO2 and CO is presented in Chapter 3,
where the data are used to interpret Spitzer spectra towards 41 low-mass objects.
1.3.2 Ice formation and grain chemistry
Ice covered grains are crucial for the interstellar chemistry leading to the formation of
complex molecules. As the embedded object starts nuclear fusion, the ices in the surrounding region are processed energetically by heating, by cosmic ray induced processes
and by ultraviolet processing. The rather simple mixtures of ices evolve to more complex ices. When the temperature reaches a high enough value the ices are desorbed and
molecules are brought back into the gas phase. Gas phase observations of this stage of star
formation indeed exhibit a large variety of complex species that originate from interstellar
7
1 Introduction
grains and thus confirm the importance of chemical reactions catalyzed on very cold ices.
Chemistry in the gas phase and on grains are thus strongly coupled.
0.
0.
0.
1.
1.
2.
O
H
C+
H2
H
H
H2O
H
H 2O
OH
Silicate grain core
H2
H2 O
CO
CO
CO
H2
N
H
H
C O
H2
H
OH
CO2 H O CH4
NH3 H2O
2
Silicate grain core
H2O-dominated ice
H2
CO
H
OH H
2.
3. Protostar
2. AV > 10
‘Late’
1. AV > 2
‘Early’
0. AV < 2
H
1.
3.
CO
H2
H2
N2
H
CO CH3 OH
CO
CO 2
CO
H2 O
CO2 CH 4
H2 O NH3 CO
Silicate grain core
CO2
CO
H2
CO
CO 2
H2 O
CO2 CH 4
H2 O NH3 CO CO2
Silicate grain core
CO-dominated ice
Figure 1.4 A proposed route of ice formation in the evolution from a diffuse to a dense
cloud. Figure is taken from Öberg [2009].
The formation of ices in the evolution from a diffuse cloud toward a protostar is illustrated in Fig. 1.4. Tielens & Hagen [1982] proposed a chemical network in which
molecules are formed from atoms which are accreted to the grain surface. In their model,
the desorption energy of the atom is larger than the energy needed for the atom to hop from
one site on the surface to the next. The atom scans the grain surface for a certain amount
of time, depending on the grain temperature and desorption energy of the atom. Meanwhile the atom may find a (radical) reaction partner on the surface, react, and form a new
species. Since hydrogen atoms are the most mobile species present on the grain, simple
H-rich species such as H2 O, CH4 , NH3 can be formed. Observations towards protostars
and dense clouds indeed point to ice layers containing H-rich molecules, dominated by
H2 O (polar ices).
Carbon monoxide is, as opposed to the other species mentioned above, efficiently
produced in the gas phase. Therefore, CO ice is formed by the freeze-out of CO directly
from the gas phase, rather than by reactions on the surface [Pontoppidan 2006]. The CO
ice forms on top of the other species in a rather pure layer, forming the so-called non-polar
ice. On the grain CO can be further processed by hydrogen addition reaction, resulting
in formaldehyde (H2 CO) and eventually methanol (CH3 OH) [Watanabe & Kouchi 2002,
Fuchs et al. 2009].
8
1.4 Mid-IR emission bands – Polycyclic Aromatic Hydrocarbons
This is not the end of the ice chemistry. Ice species, such as CH3 OH, are also subjected to VUV radiation powerful enough to photodissociate the molecules, leading to
radical species in the ice layer [e.g. Öberg et al. 2009c]. The fragments can diffuse on
the surface of the ice and react with other radical species or molecules. This system is
thought to be responsible for the formation of larger organic molecules such as methyl formate (HCOOCH3 ), formic acid (HCOOH) and (CH3 OCH3 ) [e.g. Garrod & Herbst 2006].
These species have been detected in the gas phase in regions where ices are released in
the gas phase by thermal or photo-desorption.
1.4 Mid-IR emission bands – Polycyclic Aromatic Hydrocarbons
The initial ground-based detection of the first of a family mid-IR emission features that
are now attributed to polycyclic aromatic hydrocarbons (PAHs) dates back to 1973, when
Gillett et al. discovered an unexpectedly broad emission feature peaking near 11.3 µm.
Over the next twenty years it was found that this family of bands was surprisingly widespread and associated with a wide variety of different types of astronomical objects including galactic HII regions, reflection nebulae, young stellar objects, planetary nebulae, and
post-asymptotic giant branch (AGB) objects. With the launch of ISO, and later Spitzer,
Figure 1.5 Mid-infrared spectrum of the reflection nebula NGC 7023 observed by NASA’s
Spitzer space telescope, illustrating the richness and dominance of the UIR bands. The
hatched areas are the distinct UIR bands, the shaded area are UIR plateaus. (Spectrum
from Sellgren et al. [2007], shadings courtesy Boersma)
9
1 Introduction
mapping of these features in extended objects became possible and their detection was
pushed out to galaxies across the Universe [Peeters et al. 2004b, van Dishoeck 2004,
Tielens 2008].
The most prominent of these mid-IR emission bands occur at 3.3, 6.2, 7.7, 8.6, 11.2,
and 12.7 µm and are often superimposed on broad plateaus (see Fig. 1.5). The bands
originate from regions where material is too cold to be emitting mid-IR radiation. This
requires that the carrier emits the bands upon excitation by a single photon of higher
energy (UV–near IR) and that the molecules are free gas phase species. Strong correlation
between the mid-IR emission bands and the available carbon suggests that carbon is the
main building block of the carrier. Additionally, the emission bands also originate from
regions which are dominated by harsh UV radiation, implying that the carrier must be
highly photostable. The origin of the emission features was long debated, but after more
than two decades the hypothesis that they are emitted from highly vibrationally excited
PAHs [Allamandola et al. 1989, Puget & Leger 1989] is gaining wide acceptance [e.g.
van Dishoeck 2004, Tielens 2008].
PAHs are the largest molecules known in space and contain about 10–20% of the
total available cosmic carbon. They have been found in objects, such as meteotites, and in
interplanetary dust particles, indicating their prescence in the early stages of the formation
of our solar system. PAHs may even play an important role in the formation and evolution
of life on Earth [Bernstein et al. 1999, Ehrenfreund et al. 2006].
1.4.1 The PAH building block – Carbon
Carbon is abundantly produced in stars by the triple alpha nuclear fusion process of helium, making it the sixth most abundant species in the ISM. The ability to form 4 bonds
makes carbon an important material both in a terrestrial setting as well as in space; carbon
acts as a building block from which complex organic molecules can be formed. The carbon atom contains 4 electrons which can participate in molecular bonding; two electrons
reside in the 2s atomic orbital and two electrons reside in the atomic 2p orbitals. These
atomic orbitals can mix, forming the hybridised orbitals sp, sp2 and sp3 . In the case of
the sp3 bonded form, one of the 2s electrons is promoted to the empty 2p orbital. The
2s and three 2p electron atomic electron wavefunctions mix to form sp3 atomic orbitals,
giving rise to a tetrahedral structure with the ability to form four covalent σ bonds. This
form of hybridisation is found in structures such as diamond, or in molecules such as
diamondoids (diamantane, iceane, adamantane, etc.), and alkanes (methane, ethane, etc.).
In the sp2 hybridised form only two of the three 2p orbitals mix with the 2s orbital, resulting in the ability to form three σ bonds and one π bond. This type of bonding occurs
in nanotubes, graphene, or PAHs. In the last hybridised form, sp hybridisation, the 2s
electron wavefunction only mixes with one of the p electrons. The C atom can form two
σ bonds and two π bonds. This occurs in the ethynes, such as acetylene (C2 H2 ), or carbon chain radicals (e.g. C6 H). Summarizing, carbon can reside in many different forms,
ranging from very stable configurations to highly reactive molecules. Figure 1.6 shows
some examples of the forms in which carbon atoms can be found.
10
1.4 Mid-IR emission bands – Polycyclic Aromatic Hydrocarbons
peri-condensed
cata-condensed
other carbon related species
benzene (C6 H6 )
pyrene (C16 H10 )
fullerene (C60 )
carbon chain (C6 H14 )
naphthalene (C8 H10 )
chrysene (C18 H12 )
nanotube (Cxx )
diamond (Cxx )
coronene (C24 H12 )
2,3;12,13;15,16-tribenzoterrylene (C42 H22 )
graphite (Cxx )
fragment
methylcyclohexane (C7 H14 )
Figure 1.6 Some examples of the various types of carbon containing material. (Figure
taken from Boersma [2009])
Aromatic molecules are planar structures with the atoms arranged in one or more rings
and a conjugate π-system which consists of a number of delocalized π-electrons given by
Hückels rule (4n + 2, where n = 0, 1, 2..). Polycyclic aromatic hydrocarbons, a class of
aromatic molecules, are characterized by carbon atoms arranged in chickenwire shaped
ring structures of 6 carbon atoms with 3 electrons participating in sigma-bonds and the left
over electron participating in a delocalized π-bond, resulting in a highly stable structure.
The simplest member of the stable aromatic family is benzene (C6 H6 ). The 6 carbon atom
containing hexagon of the benzene molecule forms the building block of larger aromatic
molecules consisting of 2 or more rings fused together, the PAHs. PAHs can exist in
two main forms; the peri-condensed and cata-condensed PAHs. Peri-condensed PAHs are
those which contain C atoms that are part of three fused rings of the aromatic network.
Peri-condensed PAHs are therefore centrally condensed and allow for full delocalization
of the π electron, resulting in highly stable molecules. Cata-condensed PAH molecules
do not have any carbon atoms bonded to more than two rings and therefore have a more
open structure which restricts electron delocalization making them less stable.
1.4.2 The origin of interstellar PAHs
In the ISM, PAH molecules are most likely formed in carbon-rich Asymptotic Giant
Branch (AGB) stars [Latter 1991, Cherchneff et al. 1992]. Until recently, direct evidence
for this was lacking. In general, carbon-rich AGB stars are namely too cold to efficiently
11
1 Introduction
excite the PAHs and therefore no strong PAH mid-IR emission is found towards these
objects. However, the presence of PAHs in planetary nebulae and post-AGB carbon-rich
stars, objects sampling the next stage of stellar ejecta, is unequivocal [Cerrigone et al.
2009]. Stars at this stage of their life are hotter and brighter in the near-UV and hence
pump the PAHs more efficiently, making them fluoresce in the mid-IR. Recently, Spitzer
observations of carbon-rich AGB stars have shown emission from what appears to be
a mixture of aromatic species. This mixture seems to include less stable PAH related
species that have not yet been ‘weeded out’ to the more robust PAH forms which can survive the rigors of the UV rich radiation from the hotter stars and general ISM and which
produce the well-known emission spectra.
H
C
C
H
H
H
+C H
C
2
C
H
2
H
C
C
C
H
+H
C
H
(-H )
C
H
C
+C H
2
C
C
2
H
2
(-H)
H
+H
+C H
2
H
+C H
2
2
C
C
+H
2
(-H)
C
C
H
(-H )
2
Figure 1.7 Chemical reaction scheme thought to be responsible for the production of
the first aromatic ring, from which larger PAH species grow. Figure is reproduced from
Frenklach & Feigelson [1989].
The formation process of interstellar PAHs is thought to be similar to the formation
of soot in a terrestrial setting [Frenklach & Feigelson 1989, Allamandola et al. 1989, and
Fig. 1.7]. The carbon in the outflow of carbon-rich AGB stars is mainly locked up in
CO and acetylene (C2 H2 ). Since CO is highly stable, the molecule that is most likely
responsible for the formation of PAHs in the the outflow of these stars is acetylene and its
radical derivatives. The creation of the first aromatic ring is the most problematic step in
the formation of PAHs. Hydrogen addition to a C2 H2 molecule yields the C2 H3 radical,
which can react with a second C2 H2 molecule, forming C4 H5 . Two reactions involving H
abstraction followed by reactions with two acetylene molecules yields C6 H5 , which after a
reaction of the remaining triple bond and the unpaired electron forms the first fused ring.
From here, more rings can be fused to the aromatic ring by similar acetylene addition
reactions. After their formation, they are brought into the ISM by dust-driven winds
[Speck & Barlow 1997, Boersma et al. 2006]. They can be regarded as an extension of
the grain-size distribution into the molecular (sub nanometer size) domain and are the
building blocks from which larger agglomerations — soot particles — of PAHs can be
formed.
12
1.5 Laboratory spectroscopic ice studies
1.4.3 PAHs in interstellar ices?
The mid-IR emission bands are omnipresent in space, however, the strength of these bands
varies strongly. Towards dense clouds the bands have much lower intensity. There are
two contributing factors for the quenching of the emission bands as one probes deeper
into dense clouds. The first being that the emission bands lose intensity in dense clouds
because the extinction increases and there are not enough high energy photons to excite
the PAH. The second is that the highly non-volatile PAHs condense out on grains and are
incorporated in interstellar ices.
PAHs are not expected to fluoresce in their typical mid-IR modes when incorporated
in ices, since the energy is quickly dissipated into the phonon modes of the ice lattice
[Allamandola et al. 1985, 1989]. Thus, when trapped in ices PAHs are expected to exhibit
mid-IR absorption bands instead. There are lines of evidence that support the existence of
PAHs in ices covering interstellar grains. Absorption bands likely caused by PAH feature
have been reported [Smith et al. 1989, Chiar et al. 2000, Bregman et al. 2000], but extensive laboratory studies are still lacking in the literature. Chapter 4 of this thesis describes
a study of the mid-IR spectroscopy of PAH species trapped and photolyzed in H2 O ice
with the aim to understand: 1) the roles that PAHs might play in ice processing and astrochemistry, 2) the signature PAHs add to the mid-IR spectra of embedded protostars,
and 3) identify PAH:H2 O ice photoproducts and to obtain first order estimates of their
abundances in the ices surrounding both low- and high-mass protostars.
Additional spectroscopic studies are performed in the near-UV/VIS regime on PAH
containing H2 O and CO ice in order to obtain rate constants for photoreactions of PAHs
in ices as a function of temperature. These studies are presented in Chapter 5–7. The
studies indicate that PAH are efficiently ionized and react with other ice constituent photoproducts. PAHs are thus shown to have a great impact on the interstellar ice radical
budget and charge state, particularly during the early stages of star formation and possibly also in later stages. Although much is now known about the formation of organic
molecules on interstellar ices, very little is known about the chemical processes involving
the abundantly present and largest organic molecules in the ISM, the PAHs.
1.5 Laboratory spectroscopic ice studies
Laboratory astrophysics aims to understand the physical interactions between and chemical evolution of molecular species in the interstellar medium. The physical interplay of
mixed molecular ices and their chemistry have been studied for some decades and are
reasonably well understood. The first experiments were extensions of a technique called
“Matrix Isolation Spectroscopy” [e.g. Hagen et al. 1979, 1980, Hudgins et al. 1993] and
aimed to measure band positions, FWHM and band strengths of the simplest molecular
species at cryogenic temperatures. Quickly the field evolved and more realistic “dirty
ices” — ice mixtures consisting of 2 or more species with specific mixing ratios — were
studied with the aim to understand the complex mid-IR spectra that new observatories
were discovering. Even now, these rather simple experiments still offer a wealth of infor13
1 Introduction
IR source
Sample window
Detector
Gas inlet
VUV source
Figure 1.8 A schematic of the high vacuum setup used for monitoring physical interactions and VUV induced chemical reactions in interstellar ices with mid-IR spectroscopy.
mation on the physical interactions between molecules condensed on a cold surface and
gain insight in physical parameters — such as temperature and composition — in actual
interstellar ices.
Since molecules are brought in close contact in interstellar ices, the grains act as catalytic sites for chemical reactions. These reactions are important for the overall chemistry
in the ISM. Many laboratory studies have been devoted to understanding the chemical
evolution of ices upon energetic input. Up to date, most experimental studies have employed mid-IR absorption spectroscopy on ice covered cryogenic sample windows or gold
surfaces suspended in either high- or ultrahigh vacuum systems. Recently, experimental
setups employing near-UV/VIS absorption spectroscopy have become available [Gudipati
& Allamandola 2003, and Chapter 5 of this thesis]. Both the mid-IR and near-UV/VIS
spectroscopic techniques are the subject of this thesis work and will be described shortly
in this section.
1.5.1 Mid-IR ice spectroscopy
A typical mid-IR spectroscopic setup is schematically depicted in Fig. 1.8. A sample
window is suspended in the center of a vacuum chamber, which is pumped down by a
turbomolecular pump to a pressure of 10−7 mbar. The sample window is cooled down by
a closed-cycle Helium refrigerator and the sample window temperature can be controlled
by means of resistive heating. The (mixed) gas sample is prepared off-line in a glass bulb
which can be connected to the vacuumchamber gas inlet. Ice samples are grown by vapor
depositing this gas sample onto the cold window. Subsequently, spectra are taken with
a Fourier Transform InfraRed (FTIR) spectrometer on samples of different mixing ratios
and sample window temperatures. For some of the mid-IR experiments in this thesis, energetic H2 emission is generated using a Hydrogen flow microwave (MW) discharge lamp,
to simulate energetic processing of the ices in the ISM. The resulting vacuum ultraviolet
14
1.5 Laboratory spectroscopic ice studies
(VUV) photons at 121.6 nm (Ly-α 10.2 eV) together with a broad molecular H2 emission
band at 160 nm (7.8 eV). Ices are subject to photons of high energy which may alter their
chemical identity and the chemical evolution of the photoproducts is tracked as a function
of VUV photolysis time. Typically, the FTIR spectroscopic technique has a time resolution of roughly 1 spectrum per 20 minutes for good signal to noise and a resolution of
0.5 cm−1 . Furthermore, the sample window needs to be rotated by 90◦ when changing
from the performing spectroscopy to the VUV photolysis position. Thus, this experiment
does not allow the possibility of monitoring changes in real-time nor without disturbing
the optics; requirements that must be met to fully understand the photochemistry and
determine reaction rates. This mid-IR system, however, is ideal for the identification of
functional groups in the newly formed photoproducts.
1.5.2 Near-UV/VIS absorption ice spectroscopy
Although some gas phase spectra of small PAH members are available, most of our knowledge on PAHs and related species is based on matrix isolation experiments in which the
species of interest are doped in an argon or neon matrix at low temperature, after which
the spectra — in both the mid-IR and near-UV/VIS — of the cryogenic samples are
taken. These experiments have allowed for comparison of experimental data with theoretical calculations. In addition, the experiments mentioned above using a H2 microwave
powered discharge VUV sources also allow for measuring the spectra of cationic and an-
Figure 1.9 A schematic of OASIS; the experimental setup for measuring spectroscopy and
chemical kinetics of VUV processed PAH:H2 O ice mixtures.
15
1 Introduction
ionic species. Recently, researchers realized that in the ISM PAHs should also condense
on cold grains and should be incorporated in ices. Subsequently, they can participate in
VUV induced chemical reactions and form more complex species.
The field of PAH photochemistry in realistic interstellar ice analogues was opened by
Bernstein et al. in 1999. However, it was soon realized that because PAHs have very weak
bands compared to the bands of dominant interstellar species such as, e.g., H2 O, it was
difficult to disentangle their chemistry in the laboratory with traditional IR techniques and
equally difficult to interpret the role PAHs played in the spectra of astronomical observations. The dominant interstellar ice species, however, do not have electronic transitions
and are thus largely transparent in the near-UV and visible spectral range. PAHs on the
other hand, because of their delocalized π-electrons, exhibit very strong transitions in this
part of the electromagnetic spectrum. Subsequently, an experimental setup — Optical
Absorption Setup for Ice Spectroscopy (OASIS) — aimed to study PAH electronic transitions in interstellar ice analogues was developed. A schematic the setup is displayed in
Fig. 1.9.
The new measurement technique has two major advantages compared to measurements made using mid-IR FTIR spectroscopic techniques. The first is that PAH absorptions in this wavelength regime are much stronger compared to the (very) weak PAH
absorptions in the IR (band strengths of ∼10−13 cm molecule−1 for near-UV/VIS compared to ∼10−17 cm molecule−1 for mid-IR bands). The other advantage of near-UV/VIS
studies of ices compared to IR studies is in the time resolution of the spectroscopic measurement. OASIS, on the other hand, is capable of measuring one spectrum per 5 ms. The
technique is described in more detail in Chapter 5.
1.6 Outline of this thesis
In the work presented here, two laboratory methods are employed to investigate the physical interactions and chemistry in laboratory analogues of astrophysical ices. The first
measurements are performed by FTIR studies of ices. These data are almost one-to-one
comparable to observational spectra and give good insight in the physical state of the
interstellar ice, i.e., its mixing ratio and temperature. Additionally, measurements are performed in the near-UV/VIS spectral regime. This type of spectroscopy is perfectly suited
to investigate the fast chemical reactions taking place within laboratory ice analogs of
interstellar ices with in situ VUV photolysis. This thesis is thus divided into two parts.
Part I of this thesis aims to interpret infrared laboratory measurements to explain the detection, or non-detection, of absorption bands in observational spectra. Part II aims to
qualitatively and quantitatively understand VUV driven chemical processes in PAH containing interstellar ices by means of near-UV/VIS absorption spectroscopy.
Part I: Mid-IR absorption spectroscopy
• Chapter 2 Absorption profiles and band strengths of the H2 O fundamental vibrations change in a mixed H2 O:CO ice. These changes are investigated as a function
16
1.6 Outline of this thesis
of the amount of mixed in CO. Additionally, the appearance of a CO stretching
mode band at 2152 cm−1 is quantified as a function of two physical parameters; the
amount of mixed in H2 O and the sample temperature.
• Chapter 3 The detection of NH3 ice towards low-mass protostars has long been
debated. This chapter aims to detect the NH3 umbrella mode in a set of 41 Spitzer
spectra and to derive the abundance of NH3 with respect to H2 O. Additionally,
the CH3 OH abundance is also determined from the CO stretch mode. The obtained
CH3 OH abundances are compared to previously obtained data based on the CH3 OH
ν2 C-H stretching mode.
• Chapter 4: PAHs are known to be ubiquitous in many phases of the ISM. Spectroscopy and chemistry of PAHs in H2 O ices, however, is poorly understood. This
chapter aims to obtain mid-IR spectroscopic information on PAHs trapped in H2 O
and to identify the photoproducts resulting from VUV processing of these ices. The
data are used to derive upper limits of PAH abundances in interstellar ices towards
a low- and high-mass protostar.
Part II: Near-UV/VIS absorption spectroscopy
• Chapter 5: This chapter describes a new experimental setup for performing nearUV/VIS spectroscopy on VUV processed interstellar ice analogues. The spectral
and temporal performance of the experimental setup is described by means of measurements on pyrene trapped in water ice.
• Chapter 6: The system, pyrene trapped and photolyzed in H2 O and CO ice, is
described in detail in this chapter. The chemical reactions are quantified by fitting
rate constants to the experimental data. The data are used to calculate the limit for
detecting Pyrene:H2 O ice and its photoproducts in near-UV/VIS spectra towards
dense clouds.
• Chapter 7: A set of four PAH:H2 O ice mixtures is investigated spectroscopically.
Rate constants are fitted to the experimental and a general conclusion is drawn
on the ionization of PAHs in interstellar ices. The findings are incorporated in an
astrochemical model demonstrating the importance of these processes in interstellar
environments.
• Chapter 8: This chapter is dedicated to the future prospects of the experiments
on PAH:ice spectroscopy in the Sackler Laboratory for Astrophysics and the future
prospects of the near-UV/VIS absorption spectrometer in particular. Open research
questions and possible future measurements are briefly discussed.
17
Part I
Mid-IR absorption spectroscopy
19
CHAPTER 2
Band profiles and band strengths in mixed
H2O:CO ices1
Laboratory spectroscopic research plays a key role in the identification and analysis of
interstellar ices and their structure. To date, a number of molecules have been positively
identified in interstellar ices, either as pure, mixed or layered ice structures. Previous
laboratory studies on H2 O:CO ices have employed a ‘mix and match’ principle and describe qualitatively how absorption bands behave for different physical conditions. The
aim of this study is to quantitatively characterize the absorption bands of solid CO and
H2 O, both pure and in their binary mixtures, as a function of partner concentration and
temperature. Laboratory measurements based on Fourier transform infrared transmission
spectroscopy are performed on binary mixtures of H2 O and CO ranging from 1:4 to 4:1.
A quantitative analysis of the band profiles and band strengths of H2 O in CO ice, and
vice versa, is presented and interpreted in terms of two models. The results show that a
mutual interaction takes place between the two species in the solid, which alters the band
positions and band strengths. It is found that the band strengths of the H2 O bulk stretch,
bending and libration vibrational bands decrease linearly by a factor of up to 2 when the
CO concentration is increased from 0 to 80%. By contrast, the band strength of the free
OH stretch increases linearly. The results are compared to a recently performed quantitative study on H2 O:CO2 ice mixtures. It is shown that for mixing ratios of 1:0.5 H2 O:X
and higher, the H2 O bending mode offers a good tracer to distinguish between CO2 or CO
in H2 O ice. Additionally, it is found that the band strength of the CO fundamental remains
constant when the water concentration is increased in the ice. The integrated absorbance
of the 2152 cm−1 CO feature, with respect to the total integrated CO absorption feature,
is found to be a good indicator of the degree of mixing of CO in the H2 O:CO laboratory
ice system. From the change in the H2 O absorption band strength in laboratory ices upon
mixing we conclude that astronomical water ice column densities on various lines of sight
can be underestimated by up to 25% if significant amounts of CO and CO2 are mixed in.
1 Based on: J. Bouwman, W. Ludwig, Z. Awad, K. I. Öberg, G. W. Fuchs, E. F. van Dishoeck, H. Linnartz,
Astronomy and Astrophysics, 476, 995-1003 (2007)
21
2 Band profiles and band strengths in mixed H2 O:CO ices
2.1 Introduction
Water and carbon monoxide are common constituents in vast regions of space, both in the
gas phase and in ices. Interstellar water ice was first identified in 1973 via a strong band
at 3.05 µm and unambiguously assigned to water ice following comprehensive laboratory
work [Merrill et al. 1976, Léger et al. 1979, Hagen et al. 1979]. Meanwhile, it has become
clear that H2 O ice is the most abundant ice in space. The OH stretching mode at 3.05 µm
and the H2 O bending mode at 6.0 µm are detected in many lines of sight [e.g. Willner et al.
1982, Tanaka et al. 1990, Murakawa et al. 2000, Boogert et al. 2000, Keane et al. 2001a,
Gibb et al. 2004, Knez et al. 2005] and in many different environments, ranging from quiescent dark clouds to dense star forming regions and protoplanetary disks [Whittet et al.
1988, Tanaka et al. 1994]. It has been a long-standing problem that the intensity ratio of
these two water bands in astrophysical observations is substantially different from values
derived from laboratory spectra of pure H2 O ice. In recent years it has been proposed that
this discrepancy may be due to contributions of other species, in particular more complex
organic ices, to the overall intensity of the 6 µm band [Gibb & Whittet 2002]. An alternative explanation is that the band strengths change due to interaction of H2 O molecules
with other constituents in the ice. In both high-mass and low-mass star forming regions,
CO is — together with CO2 — the most dominant species that could mix with H2 O. In
a recent study on H2 O:CO2 ices, Öberg et al. [2007a] showed indeed significant band
strength differences between pure and mixed H2 O ices. The present study extends this
work to CO containing water ice.
CO accretes onto dust grains around 20 K [Sandford et al. 1988, Acharyya et al.
2007] and plays a key role in solid state astrochemical processes, e.g., as a starting point
in hydrogenation reactions that result in the formation of formaldehyde and methanol
[Watanabe & Kouchi 2002, Hiraoka et al. 2002, Watanabe et al. 2004, Fuchs et al. 2009].
A strong absorption centered around 2139 cm−1 was assigned to solid CO by Soifer et al.
[1979], again following thorough laboratory infrared work. Further efforts in the laboratory have shown that CO molecules can be intimately mixed, either with molecules
that possess the ability to form hydrogen bonds, such as H2 O, NH3 and CH3 OH — often referred to as “polar” ices — or with molecules that can only participate in a van der
Waals type of bond, such as CO itself, CO2 and possibly N2 and O2 — so-called “nonpolar” ices. In laboratory mixtures with H2 O and CO, the two forms are distinguished
spectroscopically; the double Gaussian peak structure for the CO stretch fundamental can
be decomposed in Gaussian profiles at 4.647 µm (2152 cm−1 ) and 4.675 µm (2139 cm−1 ),
attributed to the polar and non-polar component, respectively [Sandford et al. 1988, Jenniskens et al. 1995]. On the contrary, pure CO measured in the laboratory exhibits a single
Lorentzian band, which is located around 2139 cm−1 . This Lorentzian absorption profile
can be further decomposed into three Lorentzian components centered around 2138.7,
2139.7 and 2141.5 cm−1 [H. J. Fraser, private communication].
In astronomical spectra, the 2139 cm−1 feature has been considered as an indicator
of CO in H2 O poor ice, and the 2136 cm−1 feature as CO in H2 O rich environments
[Tielens et al. 1991]. More recently it was found that the astronomical CO profiles can be
decomposed into three components at 2136.5 cm−1 , 2139.9 cm−1 and 2143.7 cm−1 , with
22
2.1 Introduction
the 2139.9 cm−1 feature ascribed to pure CO ice, and the 2143.7 cm−1 feature ascribed to
the longitudinal optical (LO) component of the vibrational transition in pure crystalline
CO [Pontoppidan et al. 2003b]. Boogert et al. [2002] proposed that the astronomically
observed peak at 2143 cm−1 can originate from CO:CO2 mixtures, but this identification
is still controversial [van Broekhuizen et al. 2006]. The assignment of the 2136.5 cm−1
feature in these phenomenological fits remains unclear. It should be noted that laboratory
and astronomical data differ slightly in peak position, largely due to the fact that grain
shape effects play a role for abundant ice molecules like CO and H2 O.
Recently, elaborate laboratory work and ab initio calculations on mixtures of CO and
H2 O have shown that the absorption around 2152 cm−1 results from CO being bound to
the dangling OH site in H2 O ice [Al-Halabi et al. 2004]. Surprisingly enough, this absorption has never been observed in the interstellar medium [e.g. Pontoppidan et al. 2003b].
The non-detection of this feature has been explained by other molecules blocking the dangling OH site, which is therefore unavailable to CO. An extension of this explanation is
that the binding sites are originally populated by CO, but that this has been processed to
other molecules, such as CO2 or methanol [Fraser et al. 2004]. Furthermore, it has been
shown that the number of dangling OH sites decreases upon ion irradiation, which in turn
results in a reduction of the integrated intensity of the 2152 cm−1 feature [Palumbo 2006,
and references therein]. The 2136–2139 cm−1 feature is ascribed to CO bound to fully
hydrogen bonded water molecules [Al-Halabi et al. 2004].
Since CO and H2 O are among the most abundant molecules in the interstellar medium,
mixed CO and H2 O ices have been subject to many experimental and theoretical studies
[e.g. Jiang et al. 1975, Hagen & Tielens 1981, Hagen et al. 1983, Al-Halabi et al. 2004,
Fraser et al. 2005]. For example, the behavior of the 2136–2139 cm−1 CO stretching
band has been quantitatively studied as a function of temperature and its band width and
position have been studied as a function of H2 O concentration in binary mixtures, but
containing only up to 25% of CO [Schmitt et al. 1989a,b]. Furthermore, water clusters
have been studied in a matrix of CO molecules with a ratio of 1:200 H2 O:CO. This has
resulted in a tentative assignment of H2 O monomers and dimers and the conclusion that
H2 O forms a bifurcated dimer structure in CO [Hagen & Tielens 1981]. Other studies
have focussed on Temperature Programmed Desorption (TPD) combined with Reflection
Absorption Infrared Spectroscopy (RAIRS) of mixed and layered CO/H2 O systems, enhancing greatly our knowledge on their structures and phase transitions [Collings et al.
2003a,b]. Nevertheless, a full quantitative and systematic study on the behavior of H2 O
in CO ice, and vice versa, with straight applications to astronomical spectra, is lacking in
the literature. This is the topic of the present work.
The desorption temperatures of CO and H2 O differ by as much as 145 K under laboratory conditions. However, H2 O/CO ices are expected to play a role in astronomical
environments at temperatures not only well below the desorption temperature of CO at
20 K [Fuchs et al. 2009], but also well above the desorption temperature of pure CO
ice, since CO can be trapped in the pores of H2 O ice [Collings et al. 2003a]. Thus far,
both species have been observed together in lines of sight. It is often concluded from
the non-detection of the 2152 cm−1 feature that H2 O and CO are not intimately mixed in
interstellar ices. On the other hand, in some lines of sight CO is trapped in pores of a host
23
2 Band profiles and band strengths in mixed H2 O:CO ices
Table 2.1 Ice mixtures and resulting deposition thicknesses used in this work. Column A
denotes the molecule of which the deposited amount is kept constant, and column B
indicates the molecule that is mixed in. The first series is used for determining the effect
of CO on the H2 O band strengths and profiles. The second series is used to determine the
effects of H2 O on the CO band strengths and profiles.
Composition
pure H2 O
pure CO
H2 O:CO 1:0.25
H2 O:CO 1:0.5
H2 O:CO 1:1
H2 O:CO 1:2
H2 O:CO 1:4
H2 O:CO 1:1
H2 O:CO 1:1
CO:H2 O 1:0.25
CO:H2 O 1:0.5
CO:H2 O 1:1
CO:H2 O 1:2
CO:H2 O 1:4
CO:H2 O 1:1
CO:H2 O 1:1
A (ML)
3000
0
3000
3000
3000
3000
3000
10000
1000
3000
3000
3000
3000
3000
10000
1000
B (ML)
0
3000
750
1500
3000
6000
12000
10000
1000
750
1500
3000
6000
12000
10000
1000
Total ice thickness (ML)
3000
3000
3750
4500
6000
9000
15000
20000
2000
3750
4500
6000
9000
15000
20000
2000
matrix, as evidenced by the detection of the 2136 cm−1 CO feature [Pontoppidan et al.
2003b]. It is plausible that this trapping results from heating of a mixture of CO and a
host molecule. Accordingly, we have also performed some experiments as a function of
temperature.
In this work, the effect of CO on the H2 O vibrational fundamentals is compared to the
effect of CO2 on these modes, as studied recently by Öberg et al. [2007a]. A comparison
between the H2 O bending mode characteristics in CO and CO2 containing ices illustrates
the sensitivity of this mode to the molecular environment. In addition, this work provides
a unique laboratory tool for investigating the amount of CO mixed with water.
The outline of this chapter is as follows. In §2.2 the experimental setup is described
and the data analysis is explained. §2.3 is dedicated to the influence of CO on the water vibrational modes, as well as the influence of water on the CO bands. In §2.4, the
astrophysical relevance is discussed and the conclusions are summarized in §2.5.
2.2 Experiment and data analysis
The experimental setup used for the measurements has been described in detail in Gerakines et al. [1995]. It consists of a high vacuum setup (≈ 10−7 Torr) in which ices are
grown on a CsI window at a temperature of 15 K. The window is cooled down by a
closed cycle He refrigerator and the sample temperature is controlled by resistive heating. A Fourier Transform InfraRed (FTIR) spectrometer is used to record ice spectra in
transmission mode from 4000 to 400 cm−1 (2.5–25 µm) with a resolution of 1 cm−1 .
24
2.2 Experiment and data analysis
The sample gas mixtures are prepared in glass bulbs, using a glass vacuum manifold.
The bulbs are filled to a total pressure of 10 mbar, which is well below the water vapor
pressure. The base pressure of the manifold is better than 10−4 mbar, resulting in negligible contamination levels. A sample of CO (Praxair 99,999%) is used without further
purification. Deionized water, further purified by three freeze-pump-thaw cycles, is used
for the H2 O:CO mixtures. Mixtures with different ratios H2 O:CO are prepared in the vacuum manifold and the resulting depositions are listed in Table 2.1. The growth rate onto
the ice is determined by setting the exposure to ∼1016 molecules cm−2 s−1 . Assuming a
monolayer surface coverage of 1015 molecules cm−2 and a sticking probability of 1, this
results in a growth rate of 10 ML·s−1 (ML=Monolayer). In the experiments where the
effect of CO on the water ice vibrational modes is investigated, the water exposure has
been kept constant with about 3000 ML of water ice for the different mixtures to facilitate
a one-on-one comparison between all samples. In the experiments where the effect of the
H2 O on the CO modes has been investigated, the total amount of deposited CO is kept
constant at 3000 ML (Table 2.1).
m
Wavelength /
3
5
10
15 20
0.16
H O
2
Absorbance / arbitrary units
0.14
CO
stretch
stretch
0.12
0.10
0.08
H O
2
0.06
free OH
H O
2
0.04
H O
2
libration
bending
0.02
0.00
4000
3500
3000
2500
2000
Wavenumber / cm
1500
1000
500
-1
Figure 2.1 A typical baseline and background corrected ice spectrum for a H2 O:CO=1:1
mixture. The measurement is performed at 15 K. The vibrational modes in the H2 O:CO
ice are indicated.
Three independent measurements are performed for 1:1 H2 O:CO mixtures. These
measurements allow for an estimate of the error in the experiment due to mixing of the
gas, deposition of the sample and other errors that may occur. A conservative error of
∼10% on the mixing ratios is deduced from these experiments. Additionally, two test
measurements are performed for samples of 1000 and 10000 ML, to check for layer thickness dependencies (Table 2.1).
The infrared spectra are taken in absorbance mode (ln(I/I0 )) using a Biorad FTS40
25
2 Band profiles and band strengths in mixed H2 O:CO ices
Table 2.2 The measured peak positions and the integration boundaries in cm−1 used to
compute the integrated intensities of the H2 O bands. The values between brackets indicate
the µm values.
Species
H2 O
CO a
a The
Assignment
νlibration
νbend
νstretch
νfree OH
ν′ stretch′
ν′ polar′
Peak
780 (12.8)
1655 (6.04)
3279 (3.05)
3655 (2.73)
2139 (4.68)
2152 (4.65)
Integration bounds
Lower
Upper
500 (20.0)
1100 (9.09)
1100 (9.09) 1900 (5.26)
3000 (3.33) 3600 (2.78)
3600 (2.78) 3730 (2.68)
2120 (4.72) 2170 (4.61)
2120 (4.72) 2170 (4.61)
integrations for the two CO bands are performed using Gaussian fits.
spectrometer. A total of 256 spectra are acquired and averaged for each sample measurement. The spectra are further processed using IDL (Interactive Data Language) in order to
flatten the baseline. This is done by fitting a second order polynomial through a set of five
points, which are visually chosen well away from absorption features. The data reduction
does not lead to a distortion of the absorption profiles. A typical absorption spectrum is
shown in Fig. 2.1.
The absorption band strengths for the three modes of pure water ice at 15 K are well
known from literature [Gerakines et al. 1995]. The adopted values are 2×10−16 , 1.2×10−17
and 3.1×10−17 cm molecule−1 for the stretching (νstretch = 3279 cm−1 or 3.05 µm), bending (νbend = 1655 cm−1 or 6.04 µm) and libration mode (νlib = 780 cm−1 or 12.8 µm),
respectively (see Fig. 2.1). Calculating the integrated absorption bands over the intervals
listed in Table 2.2 for the mixtures and scaling them to the integrated band strength for
pure water ice, allows for a deduction of the band strengths for the water ice bands in the
mixture via:
Z
Aband
H2 O
R
,
(2.1)
Aband
=
I
×
H2 O:CO=1:x
H2 O:CO=1:x
I
band
band H2 O
where Aband
HR2 O:CO=1:x is the calculated band strength for the vibrational water mode in the 1:x
mixture, band IH2 O:CO=1:x its integrated area, Aband
H2 O the band strengths for the water modes
R
as available from literature and band IH2 O the integrated area under the vibrational mode
for the pure water sample. The free OH stretching mode, “the fourth band”, is scaled to
the stretching mode for pure water since this absorption is absent in the spectrum of pure
H2 O ice.
Integration limits used throughout the experiment are listed in Table 2.2. Integrated
areas relative to the integrated area of the pure water stretching mode, A/Apure H2 O stretch, are
investigated as a function of CO concentration. For the sample with the most mixed in
CO, i.e. the 1:4 H2 O:CO mixture, an analysis in terms of cluster formation is given.
In addition, the influence of temperature on the water stretching mode is studied. The
measured spectra for the H2 O:CO mixtures are available online at the Leiden ice database.
26
W avelength /
2.7
0.05
2.8
2.75
3
6
6.5
12
7
x1/2
x1
H O:CO 1:4
m
3.2
14
16
x2
18
x2
2
0.00
H O:CO 1:2
2
0.00
H O:CO 1:1
2
0.05
0.00
H O:CO 1:0.5
2
0.05
0.00
H O:CO 1:0.25
2
0.05
0.00
pure water
free OH stretch
libration
bend
bulk stretch
0.05
0.00
3700
3650
3600
3400
3200
1700
W avenumber / cm
1600
1500
900
800
700
600
-1
27
Figure 2.2 Combined spectra of the four modes for water ice for the six measured compositions (see Table 2.1), ranging from pure
water ice (bottom) to a 1:4 H2 O:CO mixture (top figures). The spectra are taken at a temperature of 15 K. Note that the wavelength
ranges for separate modes are different. The small structures on the libration mode are experimental artifacts.
2.2 Experiment and data analysis
Absorbance / arbitrary units
0.05
2 Band profiles and band strengths in mixed H2 O:CO ices
2.3 Results
2.3.1 Influence of CO on water bands
In Fig. 2.2 the four H2 O ice fundamentals are shown for different compositions. Similar
to CO2 [Öberg et al. 2007a], CO has a clear influence on the water ice absorption bands
compared to the pure H2 O ice. This effect is different for each of the four bands. The bulk
stretch mode is most strongly affected; the band strength for this mode decreases by more
than a factor of 2 when the CO fraction is raised from 0% to 80%. The band strength of
the free OH stretch, which is absent when no CO is mixed in, is greatly enhanced with
concentration. The libration mode gradually looses intensity when the amount of CO
in the ice is increased and the peak of the absorption band shifts to lower energy. The
integrated areas of the four water modes are scaled to the pure water stretching mode and
plotted versus the CO concentration in Fig. 2.3.
[H O] /%
2
100
80
60
40
20
1.50
Error
Libration x 10
A / A
pure H O stretch
2
1.25
1.00
Stretch
0.75
Bend x 10
0.50
0.25
Free OH x 10
0.00
0
20
40
60
80
[CO] / %
Figure 2.3 The integrated intensity of the water vibration modes relative to the integrated
intensity of the pure water stretch mode plotted versus the concentration of CO in the
sample ice. It should be noted that the plots for the bending, free OH and libration mode
have been multiplied by a factor of ten to facilitate the display and that the stretch mode
is the one actually most depending on the concentration. The four water modes show to
first order a linear dependence on the CO concentration. The estimated error in the data
amounts to 10%.
A linear function Aeff = a· [CO] + b is fitted through the data points of the four water
modes. The a coefficient indicates the strength of increase/decrease of the band strength,
and the b coefficient indicates the band strength of water relative to the pure stretching
28
2.3 Results
Table 2.3 Resulting linear fit coefficients for the H2 O:CO mixtures. The coefficients indicate the strength of the interaction between CO and the H2 O host molecules in the matrix
for mixtures which are deposited at a temperature of 15 K. The corresponding values for
H2 O:CO2 ice mixtures from Öberg et al. [2007a] are listed for a comparison.
mixture
H2 O mode
H2 O:CO
νlibration
νbend
νstretch
νfree OH
νlibration
νbend
νstretch
νfree OH
H2 O:CO2
Linear Coefficients
constant (b )
slope (a )
[10−16 cm molecule−1 ] [10−19 cm molecule−1 ]
0.30±0.02
−2.1±0.4
0.13±0.02
−1.0±0.3
2.0±0.1
−16±3
0
1.2±0.1
0.32±0.02
−3.2±0.4
0.14±0.01
−0.5±0.2
2.1±0.1
−22±2
0
1.62±0.07
R2
0.93
0.84
0.95
0.99
0.99
0.81
0.99
0.99
mode when no impurities are mixed in. There exists some deviation between the data
points and the fit function which is most probably due to the deposition accuracy, but
this deviation is within the experimental error of 10%. A clear trend in all four modes
is observed. In Table 2.3, the linear fit coefficients are listed for the H2 O:CO binary
mixtures. The linear coefficient for the H2 O stretching mode is highest and negative,
indicating the strongest decrease in band strength. The free OH stretching mode has a
positive linear coefficient indicating that this is the only mode to increase in intensity
upon CO increase. A comparison with recently obtained data for H2 O:CO2 ices shows
the same trend. Apart from the bending mode, all effects are more pronounced in the CO2
mixtures, i.e., the absolute values of the a coefficient are larger by a factor of 1.3–2. This
is related to the actual interactions in the ice and work is in progress to study such effects
in more detail for a large number of species from a chemical physics perspective.
The free OH mode, the water stretching mode and the water bending mode start showing substructure superposed onto the bulk absorption profile upon increase of the fraction
of CO in the ice mixture (Fig. 2.4). The absorptions of the bulk modes are still clearly
apparent beneath the substructure. For the stretching mode this absorption shifts from
3279 cm−1 to a higher frequency of 3300 cm−1 . The peak absorption for the bulk bending
mode shifts from 1655 cm−1 for pure H2 O to 1635 cm−1 for the 1:4 H2 O:CO mixture.
The libration mode is also red-shifted upon CO concentration increase. For the pure H2 O
ice spectrum this mode is located at 780 cm−1 , while for the highest partner concentration it appears at 705 cm−1 . The free OH stretch gradually increases in frequency. The
peak absorption shifts from 3636 cm−1 for the 1:0.25 to 3655 cm−1 for the 1:4 H2 O:CO
mixture, which corresponds to a blue shift of 19 cm−1 .
The substructure which is superimposed on the bulk stretching mode around 3300
cm−1 has been previously assigned to (H2 O)n water clusters in the ice. These assignments
are based on matrix spectroscopic data of H2 O in a matrix of N2 + O2 (75:25) [Ohno
et al. 2005]. Comparable data for H2 O in a matrix of CO are not available. Neverthe29
2 Band profiles and band strengths in mixed H2 O:CO ices
less, in Fig. 2.4 it is shown that an excellent fit is obtained when peak positions from
matrix spectroscopic data are used, assuming that the peak positions are shifted which is
indicative for the difference in interaction between H2 O and CO compared to H2 O and N2
+ O2 . For each of the contributions the peak position, bandwidths (Full-Width-at-HalfMaximum (FWHM)) and integrated area are summarized in Table 2.4 and compared to
previous results obtained by Ohno et al. [2005]. Some absorptions are red-shifted and
other absorptions are blue-shifted compared to absorptions of H2 O clusters in a N2 + O2
matrix. Note that the relative H2 O concentrations in the present work are substantially
higher than in the matrix experiments by Ohno et al. [2005]. As a consequence larger
H2 O clusters are more pronounced in our spectra.
Wavelength /
2.8
3.0
m
3.2
3.4
Absorbance / arbitrary units
0.06
0.04
0.02
0.00
3800
3600
3400
3200
Wavenumber / cm
3000
2800
-1
Figure 2.4 Measured and fitted spectrum of a 1:4 H2 O:CO mixture in the frequency range
of the H2 O bulk stretch and free OH stretch modes at a temperature of 15 K. The water
stretching mode clearly shows substructure. An excellent fit is obtained by a superposition
of the bulk stretch mode and Gaussian functions representing smaller H2 O clusters in the
matrix material [Ohno et al. 2005]. The fitted peak positions and the positions from Ohno
et al. [2005] are listed in Table 2.4.
Additionally, the temperature of the sample plays a role on the band strengths. A
clear effect is encountered when the temperature is slightly increased, i.e., from 15 to
25 K, close to the desorption temperature of CO. At this temperature, the CO molecules
in the matrix start gaining enough energy to become mobile. The mobility in the matrix
allows water clusters to find partners for hydrogen bonding and to reorganize themselves
to form a stronger bulk hydrogen bonded network, as indicated by the increased bulk
stretch mode band strength and decreasing intensity of the substructure. Figure 2.5 nicely
shows the transition from H2 O clusters embedded in a matrix of CO to the formation of a
bulk hydrogen bonded network.
30
Table 2.4 Line positions, FWHMs and integrated areas of the Gaussian functions fitted to the 1:4 H2 O:CO water bulk stretch and
free OH stretch spectrum. The assignment is based on both Density Functional Theory (DFT) calculations and experimental values
from Ohno et al. [2005] obtained in a N2 /O2 matrix.
Mode
Free OH (ring) cyclic-pentamer, cyclic-trimer free OH, clusters hexamer
Asymmetric stretch H-acceptor dimer
Asymmetric stretch monomer
Symmetric stretch monomer
Symmetric stretch H-acceptor dimer
Bulk free OH stretch
Bulk free OH stretch
H-bonded OH stretch "chair" hexamer
H-bonded OH stretch "cage" hexamer
H-bonded OH stretch "prism" hexamer
H-bonded OH stretch pentamer
H-bonded OH stretch "book1, cage" hexamer
Bulk stretching mode
H-bonded OH stretch ring cyclic trimer
a Ohno et al. [2005], b This work.
Positiona
(cm−1 )
3688
3715
3715
3635
3629
—–
—–
3330
3224
3161
3368
3450
—–
3507
Positionb
(cm−1 )
3674
3693
3706
3608
3596
3655
3633
3331
3226
3175
3364
3420
3300
3506
∆ Positionb
(cm−1 )
15
22
9
27
33
—–
—–
-1
-2
-14
4
30
—–
1
FWHMb
(cm−1 )
15
7
5
5
15
19
25
43
34
40
29
86
250
42
Areab
(a.u. · cm−1 )
0.371
0.040
0.043
0.074
0.078
1.57
0.609
0.987
0.626
0.119
0.491
2.62
11.5
0.894
2.3 Results
31
2 Band profiles and band strengths in mixed H2 O:CO ices
Figure 2.6 shows the different effect of CO and CO2 on the H2 O bending mode in ice
for the 1:4 H2 O:CO2 and 1:4 H2 O:CO mixtures. The intensity ratio of the main peaks
is actually reversed in the two mixtures. The differences between the two mixtures start
showing up from a mixing ratio of 1:0.5 H2 O:X and become more pronounced for higher
CO and CO2 concentrations. In addition, the CO mixtures exhibit a stronger broad underlying feature, which is visualized by the Gaussian fit in Fig. 2.6. In other words, a detailed
study of the H2 O bending and stretching modes may provide additional information on
whether CO or CO2 dominates in the ice. The free OH stretching mode is also affected
differently by the two molecules. In the H2 O:CO2 mixtures, the mode is more shifted
to higher wavenumbers. For the 1:1 H2 O:CO2 mixture the peak position is 3661 cm−1 ,
compared to 3635 cm−1 for the H2 O:CO mixtures.
Wavelength /
2.8
m
3
3.2
0.12
30 K
Absorbance / arbitrary units
0.10
0.08
25 K
0.06
0.04
15 K
0.02
0.00
3600
3400
3200
Wavenumber / cm
3000
-1
Figure 2.5 Temperature dependence of larger clusters of water molecules in a matrix of
H2 O CO = 1:4. For increasing temperature the substructure gives way to the bulk stretch
mode when CO evaporates.
Thicker and thinner layers of the mixtures have been measured to check for thickness
dependence. We conclude that within our experimental error limit, ice thickness does not
play a significant role in the behavior of the relative band strengths. This conclusion is
supported by the observation that identical mixing ratios in the two measurement series
(H2 O/CO and CO/H2 O) show the same (scaled) spectroscopic behavior for different total
ice thicknesses (Table 2.1).
32
2.3 Results
2.3.2 Influence on the CO band
Paragraph 2.3.1 shows that mixing CO into a water mixture affects the band strengths of
the water vibrational modes. Vice versa, the CO stretch mode is also altered when water
is added to CO ice. This is seen in the experiments where the amount of deposited CO is
kept constant (Table 2.1). When water is mixed into the CO ice, the absorption changes
from a Lorentzian profile to a Gaussian profile. Furthermore, the second CO absorption
at 2152 cm−1 ascribed to CO bound to the H2 O dangling OH sites [Fraser et al. 2004]
manifests itself as a Gaussian profile and increases in integrated intensity upon increase
of water concentration. The transition from the pure CO Lorentzian shaped profile to two
Gaussian shaped profiles for the mixed CO:H2 O ices is illustrated in Fig. 2.7. We have
not further decomposed the 2139 cm−1 component.
Gaussian fits for both CO absorption components are made for the range of mixtures
as listed in Table 2.1 at a temperature of 15 K. One typical fit is shown in Fig. 2.8. With
this example it is demonstrated that excellent fits are obtained by using the fit parameters
as listed in Table 2.5. The behavior of the integrated area of the 2152 cm−1 compared with
the total integrated CO absorption is plotted as a function of CO and H2 O concentration
Wavelength /
5.8
5.9
6
m
6.1
6.2
6.3
6.4
0.025
0.020
Absorption / arbitrary units
0.015
0.010
0.005
0.000
0.025
0.020
0.015
0.010
0.005
0.000
1750
1700
1650
Wavenumber / cm
1600
1550
-1
Figure 2.6 The spectral differences in the water bending mode profile for a 1:4 H2 O:CO2
ice mixture (top) [Öberg et al. 2007a] and a 1:4 H2 O:CO mixture (bottom). The bold
(overall) spectra indicate the laboratory spectra and coincide with the fitted spectra consisting of the three Gaussian curves as indicated by the dotted lines.
33
2 Band profiles and band strengths in mixed H2 O:CO ices
in Fig. 2.9. The integrated area for the 2152 cm−1 component decreases while increasing the CO content, and is undetectable for a pure CO ice. A second order polynomial
describes how the polar component behaves with respect to the CO concentration [x] or
water concentration [100 − x] in the ice over the interval spanning from 20% CO up to a
pure CO ice. Note that the total amount of deposited CO is kept constant. The coefficients
of the second order polynomial of the form y = a · x2 + b · x + c are a = −0.005, b = 0.23
and c = 26.6. For a decreasing amount of water in the sample the peak position of the
2152 cm−1 absorption feature is most strongly affected and decreases gradually to lower
wavenumbers, until it reaches 2148 cm−1 for the 1:0.25 CO:H2 O mixture. The FWHM of
this band is also affected. It starts at a width of 10.5 cm−1 for the 1:4 CO:H2 O mixture and
decreases to 7.5 cm−1 for the 1:0.25 CO:H2 O mixture. The position of the main absorption feature at 2139 cm−1 is only slightly affected by increasing the amount of water in the
sample and decreases by 1.3 cm−1 when going from pure CO ice to the 1:4 CO:H2 O mixture, i.e., a shift toward the 2136 cm−1 feature is observed upon dilution (see Fig. 2.10).
Its position is expected to shift even more, to 2136 cm−1 for H2 O concentrations above
80%. The FWHM of this absorption feature decreases from 8 cm−1 for the 1:4 CO:H2 O
Wavelength /
4.62
0.20
4.63
4.64
4.65
4.66
m
4.67
4.68
4.69
4.7
1:4 CO:H O
2
1:2 CO:H O
Absorbance / arbitrary units
2
1:1 CO:H O
2
0.15
1:0.5 CO:H O
2
1:0.25 CO:H O
2
Pure CO
0.10
0.05
0.00
2165
2160
2155
2150
2145
2140
Wavenumber / cm
2135
2130
2125
-1
Figure 2.7 Illustration of the behavior of the CO stretch fundamental upon increase of
the concentration of H2 O in CO. The total integrated band strength remains unchanged
within the experimental error, although the maximum intensity of the absorption decreases strongly. The y-axis is cut off for the pure CO mode to make a clearer distinction
between ‘non-polar’and ‘polar’components of the CO absorption for the CO:H2 O mixtures.
34
2.3 Results
to 5 cm−1 for the 1:0.25 CO:H2 O sample. The Lorentzian peak profile of the pure CO
absorption exhibits an even smaller FWHM of 2 cm−1 . An overview of the changes in
peak position, FWHM and integrated intensity is given in Table 2.5.
Wavelength /
4.62
4.64
4.66
m
4.68
4.70
Absorbance / arbitrary units
0.15
0.10
0.05
0.00
2170
2160
2150
2140
Wavenumber / cm
2130
2120
-1
Figure 2.8 Gaussian fit of the CO stretch mode in a CO:H2 O = 1:0.25 mixture. The bold
(overall) spectrum is the measured laboratory spectrum that is reproduced by adding the
two Gaussian components (dotted lines) centered around 2138.2 cm−1 and 2147.5 cm−1 .
These are attributed to CO in a ‘non-polar’and ‘polar’environment, respectively.
Table 2.5 Lorentzian and Gaussian fit parameters for the CO stretching mode for a constant amount of CO in ice mixtures ranging from 100% CO to a 1:4 CO:H2 O mixture.
Composition
Pure CO a
1:0.25 CO:H2 O
1:0.5 CO:H2 O
1:1 CO:H2 O
1:2 CO:H2 O
1:4 CO:H2 O
a Lorentzian
Position
(cm−1 )
2138.8
2138.2
2147.5
2138.0
2148.1
2138.2
2148.3
2137.9
2149.5
2137.5
2150.0
FWHM
(cm−1 )
2.2
5.0
7.5
5.8
7.8
6.5
9.0
7.6
9.8
7.9
10.5
Area
(a.u.·cm−1 )
1.9
1.31
0.23
1.40
0.34
1.28
0.44
1.24
0.49
1.38
0.58
profile
35
2 Band profiles and band strengths in mixed H2 O:CO ices
[H O] / %
2
40
80
60
20
40
40
20
0
60
80
100
30
25
20
15
2152 cm
-1
feature / Total / %
35
10
5
0
[CO] / %
Figure 2.9 Absorbance of the 2152 cm−1 component relative to the total CO absorption as
a function of CO concentration.
The position of the 2139 cm−1 absorption feature is also strongly dependent on the
temperature of the ice as illustrated in Fig. 2.10. The majority of the CO will desorb
as the ice is heated above the CO desorption temperature. The remaining CO shows an
absorption that is shifted toward 2135 cm−1 . Thus, a shift from 2139 to 2136 cm−1 occurs
both by mixing with significant amounts of H2 O and by heating of CO:H2 O above 40 K,
even for mixtures with modest amounts of H2 O. From Fig. 2.10 it becomes clear that the
latter effect (i.e. heating) is the more critical one. One should note that the laboratory
data presented here can not be compared one-to-one with the observational data because
of shifts caused by grain shape effects.
2.4 Discussion
CO and H2 O ice abundances are often derived using the well known constants from the
literature. The experimental work presented here shows that band strengths deduced from
pure ices cannot be used to derive column densities in interstellar ices without further
knowledge on the environmental conditions in the ice. Concentrations of CO and CO2 ice
as high as 15 and 21%, respectively, relative to H2 O have been reported towards GL7009S
[Keane et al. 2001a]. If we assume a polar fraction of 75%, as reported towards W33A,
and that the polar fraction of both CO and CO2 ice are in close contact with the water, the
band strength for the H2 O bending mode is reduced by a factor of ∼1.25. In other words,
the band strength for H2 O will be smaller and hence the column density of H2 O will be
underestimated if the laboratory data from pure H2 O ice are adopted from the literature.
36
Peak postition / cm
-1
2.4 Discussion
2140.5
2140.5
2140.0
2140.0
2139.5
2139.5
2139.0
2139.0
2138.5
2138.5
2138.0
2138.0
2137.5
2137.5
2137.0
2137.0
2136.5
2136.5
2136.0
2136.0
2135.5
2135.5
0
20
40
60
0
80
20
40
60
80
100
120
140
Temperature / K
Water concentration / %
Figure 2.10 Change in peak position of the 2139 cm−1 CO absorption feature as a function
of both H2 O concentration and as a function of temperature for the 1:4 H2 O:CO mixture.
The dotted lines indicate the band positions measured for CO in the interstellar medium.
Vice versa, for deriving the column density of CO ice one should also consider the
influence of other molecules in its vicinity. The integrated absorbance of CO, however, is
less strongly influenced by the presence of H2 O. The CO absorption decreases in absolute
intensity when water is mixed in, but the total integrated band strength is compensated by
a broadening of the absorption and the appearing of the 2152 cm−1 component, at least
in the laboratory spectra. Upon heating, the CO molecules become trapped into pores as
indicated by the 2136 cm−1 feature. The absorption strength of this band is very sensitive
to temperature [Schmitt et al. 1989a]. The CO absorption profile is indicative for the
amount of water that is mixed. The percentage of CO mixed into the H2 O ice is derived in
the laboratory from the ratio between the 2152 cm−1 feature and the total CO absorption
(Fig. 2.9) in §2.3.2. This allows for a derivation of the effective band strength Aband
eff for
H2 O in a mixture with CO using the linear model proposed in § 2.3.1 and thus an estimate
of the column density for H2 O via the equation [Sandford et al. 1988]:
Ncorr =
R
τν dν
Aeff
.
(2.2)
Temperature also affects the band strength as shown in Fig. 2.5. The applications
of the model presented in Fig. 2.9 are restricted to deduction of column densities for
ices with a temperature of about 15 K and that have not been thermally processed. It
should also be noted, as recently demonstrated by Bisschop et al. [2007b], that beside
binary also tertiary mixtures have to be taken into account to compare laboratory data
with astronomical spectra.
37
2 Band profiles and band strengths in mixed H2 O:CO ices
2.5 Conclusions
Based on the experiments described in this manuscript and in recent work on H2 O:CO2
mixed ices we draw the following conclusions regarding the interaction between H2 O and
CO in a solid environment:
1. The general trend on the band strengths of the four vibrational modes in water
ices is similar for H2 O:CO and H2 O:CO2 mixtures with increasing CO or CO2
concentration. However, quantitative differences exist, reflecting differences in the
strength of the interaction, which allow to distinguish between CO and CO2 in H2 O
ice, explicitly assuming that the main constituents of the ice are H2 O, CO2 and CO.
2. The position of the water free OH stretching mode is particularly indicative of the
molecule that is interacting in the matrix (CO vs CO2 ), again under the assumption
that we only consider H2 O/CO2 /CO ices. The peak position of this mode is 26 cm−1
red-shifted for a 1:1 H2 O:CO ice mixture compared to a 1:1 H2 O:CO2 ice mixture.
3. In addition, the water bending mode is indicative of the molecule, i.e., CO or CO2 ,
that is interacting with the water ice. The relative integrated intensity of the Gaussian components reveals whether CO2 or CO is mixed into the H2 O ice. The same
restriction as mentioned in conclusions 1 and 2 applies.
4. Upon increasing the relative amount of CO in the mixture, a clear substructure
starts showing up in the bending, free OH stretch and bulk stretching mode. The
arising substructure indicates the onset of H2 O cluster formation in the H2 O:CO
ice. An assignment of the clusters has been possible following matrix isolation
spectroscopy.
5. The substructure on the stretching mode quickly gives way to the bulk water mode
when the temperature is increased close to the desorption temperature of CO. This
can easily be depicted by CO molecules becoming mobile and hence allowing single water molecules and larger water clusters to find partners for bulk hydrogen
bonding.
6. The ratio 2152 cm−1 total integrated CO absorption intensity is a tracer of the amount
of CO that is mixed into the laboratory water ice, or vice versa. In astronomical
spectra this band has not been observed.
7. H2 O column densities derived from astronomical spectra can easily be underestimated by as much as 25% when environmental influences, i.e., CO or CO2 presence, are not taken into account.
The present systematic study of CO:H2 O ice, together with recent work on CO2 :H2 O
ice, provide the tools to estimate on the mixing ratios of the three most abundant molecules
in interstellar ices.
38
CHAPTER 3
The c2d Spitzer spectroscopic survey of ices
around low-mass young stellar objects. IV. NH3
and CH3OH 1
NH3 and CH3 OH are key molecules in astrochemical networks leading to the formation
of more complex N- and O-bearing molecules, such as CH3 CN and CH3 OCH3 . Despite
a number of recent studies, little is known about their abundances in the solid state. This
is particularly the case for low-mass protostars, for which only the launch of the Spitzer
Space Telescope has permitted high sensitivity observations of the ices around these objects. In this work, we investigate the 8 − 10 µm region in the Spitzer IRS (InfraRed
Spectrograph) spectra of 41 low-mass young stellar objects (YSOs). These data are part
of a survey of interstellar ices in a sample of low-mass YSOs studied in earlier papers
in this series. We used both an empirical and a local continuum method to correct for
the contribution from the 10 µm silicate absorption in the recorded spectra. In addition,
we conducted a systematic laboratory study of NH3 - and CH3 OH-containing ices to help
interpret the astronomical spectra. We detected the NH3 ν2 umbrella mode at ∼9 µm in
low-mass YSOs for the first time. We identified this feature in 24 sources, with abundances with respect to water between ∼2 and 15%. Simultaneously, we also revisited the
case of CH3 OH ice by studying the ν4 C–O stretch mode of this molecule at ∼9.7 µm
in 16 objects, yielding abundances consistent with those derived by Boogert et al. [2008]
(hereafter paper I) based on a simultaneous 9.75 and 3.53 µm data analysis. Our study indicates that NH3 is present primarily in H2 O-rich ices, but that in some cases, such ices are
insufficient to explain the observed narrow FWHM. The laboratory data point to CH3 OH
being in an almost pure methanol ice, or mixed mainly with CO or CO2 , consistent with its
formation through hydrogenation on grains. Finally, we use our derived NH3 abundances
in combination with previously published abundances of other solid N-bearing species to
find that up to 10–20% of nitrogen is locked up in known ices.
1 Based on: S. Bottinelli, A. C. A Boogert, J. Bouwman, M. Beckwith, E. F. van Dishoeck, K I. Öberg,
K. M. Pontoppidan, H. Linnartz, G. A. Blake, N. J. Evans II and F. Lahuis, Astrophysical Journal, 718, 11001117 (2010)
39
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
3.1 Introduction
Ammonia and methanol are among the most ubiquitous and abundant (after H2 and CO)
molecules in space. Gaseous NH3 and CH3 OH are found in a variety of environments
such as infrared dark clouds, dense gas surrounding ultra-compact H II regions, massive
hot cores, hot corinos, and comets. Solid CH3 OH has been observed in the ices surrounding massive YSOs [e.g. Schutte et al. 1991, Dartois et al. 1999, Gibb et al. 2004] and
more recently toward low-mass protostars [Pontoppidan et al. 2003a]. The presence of
solid NH3 has been claimed toward massive YSOs only [Lacy et al. 1998, Dartois et al.
2002, Gibb et al. 2004, Gürtler et al. 2002], with the exception of a possible detection in
the low-mass object IRAS 03445+3242 [Gürtler et al. 2002]. However, these detections
are still controversial and ambiguous [Taban et al. 2003].
Both molecules are key participants in gas-grain chemical networks resulting in the
formation of more complex N- and O-bearing molecules, such as CH3 CN and CH3 OCH3
[e.g. Rodgers & Charnley 2001]. Moreover, UV processing of NH3 - and CH3 OH-containing ices has been proposed as a way to produce aminoacids and other complex organic
molecules [e.g. Muñoz Caro & Schutte 2003, Bernstein et al. 2002a, Öberg et al. 2009a].
In addition, the amount of NH3 in ices has a direct impact on the content of ions such
as NH+4 and OCN− , which form reactive intermediates in solid-state chemical networks.
A better knowledge of the NH3 and CH3 OH content in interstellar ices will thus help to
constrain chemical models and to gain a better understanding of the formation of more
complex, prebiotic, molecules.
During the pre-stellar phase, NH3 is known to freeze out on grains (if the core remains
starless long enough – Lee et al. 2004). Moreover, CH3 OH is known to have gas-phase
abundances with respect to H2 in hot cores/corinos that are much larger than in cold dense
clouds: ∼ (1 − 10) × 10−6 vs. ≤ 10−7 , with the former values most likely representing
evaporated ices in warm regions [e.g. Genzel et al. 1982, Blake et al. 1987, Federman et al.
1990]. Together, these findings suggest that ices are an important reservoir of NH3 and
CH3 OH and that prominent features should be seen in the absorption spectra toward highand low-mass protostars. Unfortunately, as summarized in Table 3.1, NH3 and CH3 OH
bands, with the exception of the 3.53 µm CH3 OH feature, are often blended with deep
water and/or silicate absorptions, complicating unambiguous identifications and column
density measurements. This is particularly true for NH3 whose abundance determination
based on the presence of an ammonium hydrate feature at 3.47 µm remains controversial
[e.g. Dartois & d’Hendecourt 2001].
Nonetheless, it is important to use all available constraints to accurately determine
the abundances of these two molecules. Despite the overlap with the 10 µm silicate
(Si–O stretch) feature, the NH3 ν2 umbrella mode at ∼9 µm (∼1110 cm−1 ) offers a strong
intrinsic absorption cross section and appears as the most promising feature to determine
the abundance of this species in the solid phase. Moreover, the CH3 OH ν4 C–O stretch at
∼9.7 µm (∼1030 cm−1 ) provides a good check on the validity of the different methods we
will use to subtract the 10 µm silicate absorption, since the abundance of this molecule
has been accurately determined previously from both the 3.53 and 9.75 µm features (see
Boogert et al. [2008]).
40
3.1 Introduction
Table 3.1.
Selected near- and mid-infrared features of NH3 and CH3 OH.
Mode
NH3 features:
ν3 N–H stretch
ν4 H–N–H bend
ν2 umbrella
CH3 OH features:
ν2 C–H stretch
ν6 & ν3 –CH3 deformation
ν7 –CH3 rock
ν4 C–O stretch
Torsion
λ (µm)
ν̄ (cm−1 )
2.96
6.16
3375
1624
9.00
1110
Blended with H2 O (O–H stretch, 3.05 µm/3275 cm−1 )
Blended
with
H2 O
(H–O–H
bend,
5.99 µm/1670 cm−1 ), HCOOH
Blended with silicate
3.53
6.85
8.87
9.75
14.39
2827
1460
1128
1026
695
–
Blended (e.g. with NH+4 )
Weak; blended with silicate
Blended with silicate
Blended with H2 O libration mode
Problem
Note. — The bold-faced lines indicate the features studied here.
Note. — The nomenclature for the NH3 and CH3 OH vibrational modes are adopted from Herzberg [1945].
More detailed spectroscopic information is particularly interesting for low-mass protostars as the ice composition reflects the conditions during the formation of Sun-like
stars. Such detections have only become possible with Spitzer, whose sensitivity is necessary to observe low luminosity objects even in the nearest star-forming clouds.
The gain in sensitivity offered by Spitzer compared to previous space-based observatory, as well as the spectral resolution of the data analyzed here (∆λ/λ ∼ 100), imply
that the interpretation of the astronomical spectra should be supported by a systematic
laboratory study of interstellar ice analogues containing NH3 and CH3 OH. The spectral
appearance of ice absorption features, such as band shape, band position and integrated
band strength, is rather sensitive to the molecular environment. Changes in the lattice
geometry and physical conditions of an ice are directly reflected by variations in these
spectral properties. In the laboratory, it is possible to record dependencies over a wide
range of astrophysically relevant parameters, most obviously ice composition, mixing ratios, and temperature. Such laboratory data exist for pure and some H2 O-rich NH3 - and
CH3 OH-containing ices [e.g. D’Hendecourt & Allamandola 1986, Hudgins et al. 1993,
Kerkhof et al. 1999, Taban et al. 2003], but a systematic study and comparison with observational spectra is lacking.
In principle, the molecular environment also provides information on the formation
pathway of the molecule. For example, NH3 ice is expected to form simultaneously with
H2 O and CH4 ice in the early, low-density molecular cloud phase from hydrogenation
of N atoms [e.g. Tielens & Hagen 1982]. In contrast, solid CH3 OH is thought to result
41
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
primarily from hydrogenation of solid CO, a process which has been confirmed to be
rapid at low temperatures in several laboratory experiments [e.g. Watanabe & Kouchi
2002, Hidaka et al. 2004, Fuchs et al. 2009]. A separate, water-poor layer of CO ice is
often found on top of the water-rich ice layer in low-mass star-forming regions due to
the ‘catastrophic’ freeze-out of gas-phase CO at high densities [Pontoppidan et al. 2003a,
Pontoppidan 2006]. Hydrogenation of this CO layer should lead to a nearly pure CH3 OH
ice layer [e.g. Cuppen et al. 2009], which will have a different spectroscopic signature
from that of CH3 OH embedded in a water-rich matrix. The latter signature would be
expected if CH3 OH ice were formed by hydrogenation of CO in a water-rich environment
or by photoprocessing of H2 O:CO ice mixtures, another proposed route [e.g. Moore &
Hudson 1998].
Here, we present Spitzer spectra between 5 and 35 µm of ices surrounding 41 lowmass protostars, focusing on the 8 − 10 µm region that contains the ν2 umbrella and ν4
C–O stretch modes of NH3 and CH3 OH, respectively. This chapter is part of a series of
ice studies [Boogert et al. 2008, Pontoppidan et al. 2008, Öberg et al. 2008] carried out
in the context of the Spitzer Legacy Program “From Molecular Cores to Planet-Forming
Disks” (“c2d”; Evans et al. 2003). In §3.2, we carry out the analysis of the Spitzer data
in 8 − 10 µm range. In §3.3, we present the laboratory data specifically obtained to
help interpret the data that are discussed in §3.4. Finally, we conclude in §3.5 with a
short discussion of the joint astronomy-laboratory work (including the overall continuum
determination).
3.2 Astronomical observations and analysis
The source sample consists of 41 low-mass YSOs that were selected based on the presence of ice absorption features. The entire sample spans a wide range of spectral indices
α = −0.25 to +2.70, with α defined as d log(λFλ )/d log(λ), where d indicates the derivative, and Fλ represents all the photometric fluxes available between λ = 2.17 µm (2MASS
Ks -band) and λ = 24 µm (Spitzer/MIPS band). In the infrared broad-band classification
scheme, 35 out of 41 objects fall in the embedded Class 0/I category (α > 0.3). The
remaining 6 objects are flat-spectrum type objects [−0.3 < α < 0.3; Greene et al. 1994].
Spitzer/IRS spectra (5-35 µm) were obtained as part of the c2d Legacy program (PIDs 172
and 179), as well as a dedicated open time program (PID 20604), and several previously
published GTO spectra [Watson et al. 2004]. We refer the reader to Table 1 and Section 3
of Boogert et al. [2008] for the source coordinates and a description of the data reduction
process (including overall continuum determination).
As mentioned previously, spectral signatures in the 8−10 µm region are dominated by
the Si–O stretching mode of silicates. The overall shape as well as the sub-structure of the
silicate feature depend on grain size, mineralogy, level of crystallinity. These effects are
degenerate and so these different factors cannot be easily separated. For example, large
grains and the presence of SiC both produce a shoulder at 11.2 µm [e.g. Min et al. 2007].
Therefore, trying to fit the 10 µm silicate feature by determining the composition and
42
3.2 Astronomical observations and analysis
size of the grains is a complex process. For this reason, we use two alternative methods to
model the silicate profile and extract the NH3 (and CH3 OH) feature(s) from the underlying
silicate absorption.
3.2.1 Local continuum
The first method uses a local continuum to fit the shape of the silicate absorption. For
this, we fit a fourth order polynomial over the wavelength ranges 8.25–8.75, 9.23–9.37,
and 9.98–10.4 µm, avoiding the positions where NH3 and CH3 OH absorb around 9 and
9.7 µm. These fits are shown with thick black lines in Fig. 3.1. After subtraction of the
local continuum from the observations, we fit a Gaussian to the remaining NH3 and/or
CH3 OH feature, when present, as shown in Fig. 3.2. The results of the Gaussian fits are
listed in Table 3.5 of Appendix 3.6.
3.2.2 Template
The second method assumes that the 8–10 µm continuum can be represented by a template silicate absorption feature, selected among the observed sources. A comparison of
the results obtained using a template to those obtained using a simple local continuum
provides an estimate of the influence of the continuum choice on the shape and depth of
the NH3 and CH3 OH features. The templates were chosen using an empirical method.
Upon examination of the 10 µm feature of the entire sample, the sources could be separated into three general categories, depending on the shape of the wing of the silicate
absorption between ∼8 and 8.7 µm (which we will refer to as the 8 µm wing): (i) sources
with a straight 8 µm wing (Fig. 3.3-a), (ii) sources with a curved 8 µm wing (Fig. 3.3-b),
and (iii) sources with a rising 8 µm wing (“emission” sources, Fig. 3.3-c).
Note that, since radiative transfer in the 8–10 µm region can be complicated by the
presence of silicate emission, we only consider sources that are the least affected by emission, that is those falling in one of the first two categories. Nevertheless, non-rising silicate
profiles might still suffer from the presence of emission. To try and estimate the impact
of this potential effect, we used two silicate emission sources from Kessler-Silacci et al.
[2006], and subtracted these emission profiles from our absorption profiles, assuming that
the emission represented 10 to 50% of the observed absorption. After removal of a local
continuum, we determined the integrated optical depths of the NH3 and CH3 OH features
in the spectra corrected for emission, and compared these to the integrated optical depths
of the uncorrected spectra. We find that the difference can be up to a factor of two and
therefore identify this possible presence of underlying emission as the largest source of
uncertainty in our abundance determinations.
For each of the straight and curved 8 µm wings, two sources (in order to test for
template-dependent effects) were selected as possible templates for the silicate feature.
The selection criteria were: (i) a silicate feature as deep as possible to minimize the effects
43
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Figure 3.1 (Top) Local continuum (thick blue/black lines) and template (red/grey lines) fits
to all sources for which a template could be found. (See §3.2.2 for details) — (Bottom)
Local continuum fits to emission sources or sources for which no reasonable template
could be found.
44
3.2 Astronomical observations and analysis
Figure 3.2 (Top) Residual after removal of local continuum and template fits for all sources
for which a template could be found. (See §3.2.2 for details) — (Bottom) Residual after
removal of local continuum fits for emission sources or sources for which no reasonable
template could be found.
45
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Figure 3.3 Examples illustrating the three shapes of the 8 µm wing shown by the thick
grey line: (a) straight, (b) curved, and (c) rising.
of silicate emission and (ii) little NH3 and CH3 OH signal, as estimated after subtraction
of a local continuum. Additionally, we added to this list the GCS3 spectrum observed by
Kemper et al. [2004] toward the Galactic Center. The spectra of these templates in the
8–10 µm region are displayed in Fig. 3.4.
For all the other sources in our sample, the best template was determined by scaling
the possible templates to the observed optical depth at different wavelengths (8.75, 9.30,
9.37, 9.70, 9.98 µm) and finding the combination (template + scaling point) that gave
the least residuals over the same wavelength ranges used to estimate the local continuum
(8.25–8.75, 9.23–9.37, 9.98–10.4 µm). The result of this process is displayed for each
source in the top part of Fig. 3.1, where the best template is shown by a grey line. The
bottom panels of Fig. 3.1 show sources for which no reasonable template could be found,
as well as emission sources, in which case only the local continuum is overlaid. As in the
case of the local continuum method, the spectra obtained after subtraction of the templates
are shown in Fig. 3.2. Taken together, NH3 features are detected in 24 out of 41 sources.
The top panel of Figure 3.2 shows that the CH3 OH feature is not affected by the continuum choice, whereas the width of the NH3 band is somewhat sensitive to this choice,
especially if there is no CH3 OH absorption, in which case the local continuum yields
a wider NH3 profile. For both continua, there is clearly a feature around 9 µm, which
we attribute to NH3 , with the characteristics and limitations given and discussed in the
following sections.
3.2.3 NH3 ice column densities and abundances
Gaussian fits were performed to the NH3 and/or CH3 OH features when present, and derived parameters for NH3 are listed in Table 3.5 (Appendix 3.6). Table 3.2 gives the
column densities derived for NH3 for each of the two methods employed to determine the
continuum, using a band strength of 1.3×10−17 cm molecule−1 for the NH3 ν2 umbrella
mode appropriate for a water-rich ice [D’Hendecourt & Allamandola 1986, Kerkhof et al.
1999]. The two methods generally agree to within a factor of 2 or better. A similar factor
of ≤2 overall uncertainty is estimated for those sources for which only the local continuum
has been used.
46
3.2 Astronomical observations and analysis
The position of the NH3 ν2 umbrella mode is very close to that of the ν7 CH3 -rock
mode of CH3 OH. As illustrated by our laboratory data (see §3.3), sources with an absorption depth at ∼9.7 µm (CO-stretch mode of CH3 OH) at least twice as large as the absorption depth at ∼9 µm (blend of CH3 -rock mode of CH3 OH and NH3 umbrella mode) have
a significant contribution to the 9 µm integrated optical depth from the CH3 -rock mode
of CH3 OH. In these cases (sources followed by an asterisk in Table 3.2 and in Table 3.5
of Appendix 3.6), we performed the following correction: we scaled a H2 O:CH3 OH=9:1
laboratory spectrum to the observed optical depth of the CO-stretch mode of CH3 OH,
determined the integrated optical depth of the CH3 -rock mode of CH3 OH in that scaled
spectrum, and subtracted it from the total observed optical depth at 9 µm. This correction is justified by the fact that the H2 O:CH3 OH:NH3 =10:4:1 spectrum, a typical interstellar abundance mixture, is well reproduced around 8–10 µm by a combination of
H2 O:CH3 OH=9:1 and H2 O:NH3 =9:1 (see § 3.3).
Figure 3.4 Silicate features of the sources used as templates for a straight 8 µm wing
(left), curved 8-µm wing (middle), and GCS3 (right). The bottom panels of each plot are
the residuals after removal of the local continuum shown in grey in the top panels. The
optical depth scale is kept fixed for comparison. These sources are selected to have no or
at most weak NH3 and CH3 OH absorptions.
47
Source
IRAS 03235+3004
L1455 IRS3
IRAS 03254+3050
B1-b∗
IRAS 04108+2803
HH 300
IRAS 08242−5050
IRAS 15398−3359
B59 YSO5
2MASSJ17112317−272431
SVS 4-5∗
R CrA IRS 5
RNO 15c
IRAS 03271+3013
B1-a
L1489 IRS
IRAS 13546−3941
RNO 91
IRAS 17081−2721
EC 74c
EC 82
EC 90
EC 92∗
CrA IRS7 B∗
L1014 IRS
CK4
NH3 column densitiesa and abundances with respect to H2 O iceb
NH3 , local
×1017 cm−2
% H2 Ob
6.83 ( 0.98)
0.57 ( 0.23)
2.44 ( 0.39)
∼7.3
1.23 ( 0.24)
0.90 ( 0.22)
4.77 ( 0.46)
8.73 ( 1.18)
4.92 ( 0.72)
13.10 ( 1.06)
∼2.4
0.91 ( 0.23)
0.80 ( 0.21)
4.90 ( 0.88)
3.46 ( 0.69)
2.31 ( 0.30)
0.94 ( 0.16)
2.03 ( 0.30)
0.86 ( 0.16)
1.00 ( 0.29)
1.22 ( 0.14)
0.67 ( 0.20)
∼0.5
∼3.0
3.72 ( 0.91)
0.84 ( 0.13)
4.71 ( 1.00)
6.21 ( 3.51)
6.66 ( 1.37)
∼4.2
4.29 ( 1.03)
3.46 ( 0.90)
6.13 ( 0.85)
5.90 ( 1.77)
3.53 ( 0.88)
6.70 ( 0.54)
∼4.3
2.54 ( 0.67)
11.58 ( 3.18)
6.37 ( 1.86)
3.33 ( 0.98)
5.42 ( 0.96)
4.56 ( 0.87)
4.78 ( 0.81)
6.54 ( 1.39)
9.35 ( 3.13)
31.31 ( 6.65)
3.94 ( 1.24)
∼3.0
∼2.8
5.20 ( 1.43)
5.37 ( 0.86)
NH3 , template
×1017 cm−2
% H2 Ob
8.94 ( 1.03)
1.41 ( 0.27)
4.58 ( 0.49)
∼9.8
2.07 ( 0.39)
2.23 ( 0.37)
4.41 ( 0.54)
13.80 ( 1.35)
6.37 ( 0.99)
20.60 ( 2.76)
∼5.8
1.49 ( 0.31)
–
–
–
–
–
–
–
–
–
–
–
–
–
–
6.17 ( 1.20)
15.37 ( 6.86)
12.52 ( 2.10)
∼5.6
7.21 ( 1.69)
8.60 ( 1.65)
5.66 ( 0.89)
9.33 ( 2.65)
4.57 ( 1.17)
10.58 ( 1.42)
∼10.3
4.15 ( 0.92)
–
–
–
–
–
–
–
–
–
–
–
–
–
–
Template
IRAS 12553
GCS3
IRAS 12553
IRAS 12553
IRAS 23238
DG Tau B
IRAS 12553
IRAS 12553
CrA IRS7 A
IRAS 23238
GCS3
IRAS 12553
–
–
–
–
–
–
–
–
–
–
–
–
–
–
Scaling point
µm
9.30
9.37
10.40
9.70
9.70
9.70
9.70
9.70
9.70
9.70
8.75
9.70
–
–
–
–
–
–
–
–
–
–
–
–
–
–
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
48
Table 3.2.
Table 3.2.
Source
0.20
17.28
15.10
0.24
11.93
0.31
0.47
0.47
0.61
0.28
0.52
0.46
0.97
5.44
1.60
4.15
4.40
8.29
5.93
4.04
3.10
2.09
2.05
2.04
0.93
1.23
1.89
0.89
10.35
1.24
NH3 , template
×1017 cm−2
% H2 Ob
3-σ upper limits
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
Template
Scaling point
µm
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
–
Note. — Sources in bold were used as templates. Uncertainties quoted in parenthesis are statistical errors from the
Gaussian fits while absolute errors are up to a factor of 2.
a
Derived using a band strength of 1.3×10−17 cm molecule−1 .
b
Using the H2 O ice column densities listed in Paper I.
c
Values are likely upper limits (see §3.4.2 for details).
∗
Sources with τ9.7µm > 2 × τ9.0µm , for which an estimated contribution from the CH3 -rock mode of CH3 OH was
subtracted (see text for details).
49
3.2 Astronomical observations and analysis
LDN 1448 IRS1
IRAS 03245+3002
L1455 SMM1
IRAS 03301+3111
B1-c
IRAS 03439+3233
IRAS 03445+3242
DG Tau B
IRAS 12553-7651
Elias 29
CRBR 2422.8−342
HH 100 IRS
CrA IRS7 A
CrA IRAS32
IRAS 23238+7401
NH3 , local
×1017 cm−2
% H2 Ob
Cont’d
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
The inferred NH3 ice abundances range from . 1% to 15% with respect to H2 O ice,
excluding the abnormally high value of EC 82. When considering all values (except
that of EC 82) determined with the local continuum method, this relative abundance is
centered on 5.3% with a standard deviation of 2.0%. If we use values determined with the
template method whenever available, we find a mean of 7.0±3.2%. Either way, within the
errors, this is similar to what was obtained by Öberg et al. [2008] for CH4 (4.7±1.6%),
another ice component that should form via hydrogenation. For 6 out of the 8 sources
where both NH3 and CH4 are detected, the NH3 -to-CH4 ratio is slightly larger than 1
(∼1.2). Based on elemental abundance ratios, one would expect NH3 /CH4 smaller than
1, but since two thirds of the carbon is in refractory grains and some fraction of the
gaseous CO locked up in CO at the ice formation threshold, NH3 -to-CH4 ratios larger
than 1 are consistent with both NH3 and CH4 being formed by hydrogenation of N and C,
respectively.
Here, we only report values for the Gaussian parameters and derived column densities
in the appendix (see Table 3.5), to show that the numbers we obtain in this independent study are consistent with those reported in Paper I. Our recommended abundances
are those from paper I, based on the combined 9.75 and 3.53 µm analysis. The inferred
CH3 OH abundances range from < 1% to > 25% with respect to H2 O ice, indicating significant CH3 OH/NH3 abundance variations from source-to-source. Such relative abundance
variations can already be clearly seen from the changing relative depths of the 9.0 and
9.7 µm features (see also Paper I). Thus, NH3 and CH3 OH ice are likely formed through
different formation pathways and/or in different ice environments.
3.3 Laboratory work and analysis
The band profiles presented in Fig. 3.2 contain information on the ice environment in
which NH3 and CH3 OH are located, and thus their formation and processing history. To
extract this information, a systematic laboratory study of the NH3 and CH3 OH features
in a variety of ices has been carried out. Specifically, three features between 8 and 10 µm
have been analyzed:
1. the NH3 ν2 umbrella mode, at ∼9.35 µm or 1070 cm−1 in pure NH3 ice, and
with band strength Apure =1.7×10−17 cm molecule−1 [D’Hendecourt & Allamandola 1986],
2. the CH3 OH ν4 CO– stretching mode, at ∼9.74 µm or 1027 cm−1 in pure CH3 OH
ice, and with Apure = 1.8×10−17 cm molecule−1 [D’Hendecourt & Allamandola
1986],
3. the CH3 OH ν7 CH3 rocking mode, at ∼8.87 µm or 1128 cm−1 in pure CH3 OH ice,
and with Apure = 1.8×10−18 cm molecule−1 [Hudgins et al. 1993].
It should be noted that, as mentioned in the above list, the quoted positions are for pure
ices only and therefore slightly deviate from the astronomical values given in Table 3.1.
50
3.3 Laboratory work and analysis
This laboratory study targeted pure, binary and tertiary interstellar ice analogs consisting of different mixtures of H2 O, NH3 , CH3 OH, CO and CO2 , the major ice components.
All measurements were performed under high vacuum conditions (∼ 10−7 mbar) using an
experimental approach described in Gerakines et al. [1995], Chapter 2 of this thesis, and
Öberg et al. [2007a]. The ice spectra were recorded in transmission using a Fourier transform infrared spectrometer covering 25–2.5 µm (400–4000 cm−1 ) with 1 cm−1 resolution and by sampling relatively thick ices, typically several thousands monolayers (ML)1
thick. These ices were grown at a speed of ∼1016 molecules cm−2 s−1 (10 MLs−1 on a
temperature-controlled CsI window.
Figure 3.5 Example of a reduced laboratory spectrum (solid black line) for a
H2 O:CH3 OH:NH3 = 10:4:1 ice mixture at 15 K, in the 8–10 µm / 960–1220 cm−1
range. This spectrum can be approximated as the sum (solid green/dark grey line) of
H2 O:CH3 OH=9:1 (solid red/light grey line) and H2 O:NH3 =9:1 (dash-dot blue/grey line).
The bottom plot is the difference between the two, showing that the feature at 9 µm (blend
of NH3 and CH3 OH CH3 -rock modes) is well reproduced by the sum of the two individual
signatures. This figure also illustrates the fact that the positions of the features in mixed
ices differ from that in pure ices (see list at the beginning of this section).
A typical reduced spectrum for an ice mixture containing H2 O:CH3 OH:NH3 = 10:4:1
at 15 K is shown in Fig. 3.5. Since band profiles and strengths change with ice composition and also with temperature, the three fundamentals mentioned above were investigated
as a function of temperature ranging from 15 to 140 K with regular temperature steps for
1 One ML corresponds to the layer thickness resulting from an exposure for 1 second at a pressure of 10−6
torr assuming a sticking probability of one. One ML is equivalent to about 1015 molecules cm−2 .
51
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
a number of binary and tertiary mixtures (listed in Appendix 3.6). An IDL routine was
used to determine the location of the band maximum, FWHM and integrated absorbance
of the individual absorption bands. For the asymmetric NH3 ν2 umbrella mode the band
position has been determined by the maximum absorbance and for the symmetric profiles
the spectral parameters have been determined from Gaussian fits of baseline subtracted
spectra. The resulting absolute frequency uncertainty is of the order of 1 cm−1 . The
measurements are presented in Table 3.6 of Appendix B, and are included in the Leiden
laboratory database2 .
Figure 3.6 (Left) FTIR ice spectra of the νNH3 mode for pure NH3 , a H2 O:NH3 =1:1
and a H2 O:NH3 =9:1 mixture at a temperature of 15 K. At the low frequency side of the
spectrum the H2 O libration mode (centered around 770 cm−1 , or 13 µm) starts showing
up for the H2 O-containing mixtures. — (Right) Temperature effect on a H2 O:NH3 =9:1
mixture: decreasing FWHM with increasing temperature.
NH3 and CH3 OH both have the ability to form hydrogen bonds in water-rich matrices,
so it is not surprising that the band profile changes compared with pure ices because of the
various molecular interactions [e.g., D’Hendecourt & Allamandola 1986]. In addition to
profiles, band strengths can change with environment and with temperature, as discussed
for the cases of CO and CO2 in water-rich ices in Kerkhof et al. [1999], Öberg et al.
[2007a], and Chapter 2 of this thesis. Figure 3.6 shows how the NH3 ν2 umbrella mode
absorption maximum shifts from 1070 cm−1 (9.35 µm) for pure NH3 ice to 1118 cm−1
(8.94 µm) for an astronomically more realistic H2 O:NH3 =9:1 (hereafter 9:1) mixture, for
which the FWHM and integrated band strength also change significantly. For example,
the band strength is lowered in the 9:1 mixture to 70% of its initial value in pure NH3 ice.
This is in good agreement with previous experiments performed by Kerkhof et al. [1999].
The spectral appearance also depends on temperature; for the 9:1 mixture a temperature
increase from 15 to 120 K results in a redshift of the peak position from 1118 to 1112 cm−1
(8.94 to 8.99 µm) and the FWHM decreases from 62 to 52 cm−1 (0.50 to 0.42 µm) (see
Fig. 3.7). The NH3 band strength, on the other hand, does not show any temperature
dependence.
2 http://www.strw.leidenuniv.nl/∼lab/
52
3.3 Laboratory work and analysis
Figure 3.7 A plot indicating the changes in peak position (left) and FWHM (right) of the
NH3 ν2 umbrella mode as a function of temperature in a 9:1 H2 O:NH3 ice.
If NH3 is in a water-poor environment with CO and/or CO2 , the ν2 peak position
shifts to the red compared with pure NH3 , to as much as 1062 cm−1 (9.41 µm). The
FWHM is not much affected whereas the band strength is lowered by 20%. Because of
the intrinsically large difference in band maximum position between NH3 in a water-poor
and water-rich environment, the astronomical observations can distinguish between these
two scenarios.
Methanol-containing ices have been studied in a similar way (see Fig. 3.8). The
weakly absorbing ν7 CH3 rocking mode at ∼1125 cm−1 (8.89 µm) is rather insensitive
Figure 3.8 (Left) Spectra of the CH3 OH νCO modes and νCH3 modes for pure CH3 OH, a
H2 O:CH3 OH=1:1, a H2 O:CH3 OH=9:1 and a CO:CH3 OH=1:1 ice mixture at a temperature of 15 K. — (Right) Temperature effect on the CO-stretch mode of a H2 O:CH3 OH=9:1
mixture.
53
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
to H2 O mixing, but the ν4 CO stretch vibration shifts to the red from 1028 to 1020 cm−1
(9.73 to 9.80 µm) when changing from a pure CH3 OH ice to a H2 O:CH3 OH=9:1 mixture.
In the latter spectrum the CH3 OH ν4 CO stretch mode needs to be fitted with a double
Gaussians. A substructure appears for a temperature of 80 K (right panel of Fig. 3.8)
while for even higher temperatures, a clearly double peaked structure becomes visible (as
previously seen in e.g. Fig. 2 of Schutte et al. 1991). This splitting hints at different physical sites and has been previously ascribed to type II clathrate formation in the ice [Blake
et al. 1991].
Figure 3.9 Spectra of CH3 OH:CO mixtures in the range of the methanol CO stretch mode
and the methanol CH3 rock mode. A small blue shift together with a clear substructure
are seen upon mixing in more CO.
When CH3 OH is mixed with CO, the band maximum shifts from 1028 to 1034 cm−1
(9.73 to 9.67 µm) when going from a 9:1 to a 1:9 CH3 OH:CO mixture. When 50% or
more CO is mixed in, the CH3 OH ν4 CO stretch mode starts to show a shoulder and cannot
be fitted correctly by a single Gaussian component (see Fig. 3.9). Such a two-component
profile would not be recognized, however, at the spectral resolution and signal/noise of
our Spitzer data, so for the comparison between laboratory and observational data a single
Gaussian is used. Overall, the shifts of the CH3 OH ν4 mode between water-rich and COrich mixtures are much smaller than in the case of the NH3 ν2 mode.
The effect of CH3 OH on the 4.7µm ν1 stretch mode of CO has also been investigated.
The band maximum shifts from 2139 cm−1 (4.68 µm) for the nearly pure 9:1 CO:CH3 OH
mixture to 2136 and 2135 cm−1 for the 1:1 and 1:9 mixtures, respectively. The CO band
located at 2136 cm−1 is often referred to as CO residing in a polar, mainly H2 O ice,
environment. Clearly, the polar CH3 OH molecules can also contribute to CO absorption
at 2136 cm−1 when intimately mixed in an astronomical ice.
Binary mixtures of NH3 and CH3 OH have been studied as well. The CH3 OH modes
behave very much as they do in a pure methanol ice, but the NH3 ν2 umbrella mode is
54
3.3 Laboratory work and analysis
clearly suppressed. Its integrated absorbance is readily reduced to 70% of the integrated
absorbance of pure NH3 in a CH3 OH:NH3 =1:1 mixture and becomes even lower for a 4:1
binary composition. The NH3 band also broadens compared to pure NH3 or H2 O:NH3
mixtures and strongly overlaps with the CO stretching mode of CH3 OH, to the level that
it becomes difficult to measure.
A qualitative comparison with the astronomical data (see §3.4) indicates that neither
pure NH3 , CH3 OH, nor mixed CH3 OH:NH3 or H2 O-diluted binary ices can simultaneously explain the different NH3 profiles in the recorded Spitzer spectra. Thus, a series
of tertiary mixtures with H2 O:CH3 OH:NH3 in ratios 10:4:1, 10:1:1 and 10:0.25:1 have
been measured, because CH3 OH is the next major ice component. These ratios roughly
span the range of observed interstellar column density ratios. In Fig. 3.10, the spectra
of H2 O:CH3 OH:NH3 tertiary mixtures are plotted and compared to binary H2 O:CH3 OH
and H2 O:NH3 data. The NH3 ν2 umbrella mode shifts slightly to the blue in the presence of both H2 O and CH3 OH, with an absorption maximum at 1125 cm−1 (8.90 µm)
for the 10:4:1 H2 O:CH3 OH:NH3 mixture (compared to 1118 cm−1 (8.94 µm) in the
H2 O:NH3 =9:1 mixture). The peak intensity of the NH3 ν2 umbrella mode band in this tertiary mixture is small compared with that of the CH3 OH CH3 rock mode, but its integrated
intensity is a factor of two larger because of the larger NH3 width.
Figure 3.10 Normalized spectra of the CH3 OH ν4 C–O mode (right panel),
and NH3 ν2 umbrella mode (left panel) for a H2 O:CH3 OH:NH3 =10:0.25:1, a
H2 O:CH3 OH:NH3 =10:1:1 and a H2 O:CH3 OH:NH3 =10:4:1 mixture at a temperature of
15 K. These mixture ratios span the range of observed interstellar column density ratios. Spectra were normalized to better show the changes in band maximum position
and FWHM of each feature. Spectra of a H2 O:CH3 OH=9:1 and a H2 O:NH3 =1:1 mixture were offset and overlaid in light grey in the right and left panel, respectively. In the
case of H2 O:CH3 OH:NH3 =10:4:1, the NH3 ν2 umbrella mode is heavily blended with the
CH3 OH ν7 CH3 rocking mode, so that the dark grey line actually shows the Gaussian fit
to the underlying NH3 feature, whereas the full 9 µm feature is shown in black.
55
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
The ν4 C–O stretching vibration profile of CH3 OH in the tertiary mixture does not
differ much from the binary values for the highest water content. The position of the
absorption maximum is also only marginally affected by the temperature. The FWHM
decreases from 30 cm−1 (0.29 µm) for the 10:4:1 mixture to 22 cm−1 (0.21 µm) for the
10:0.25:1 mixture.
Besides H2 O, other species may also be regarded as potential candidates for changing the spectral appearance of the NH3 and/or CH3 OH features. Chemically linked is
HCOOH [Bisschop et al. 2007a] which unfortunately cannot be deposited in the present
setup because of its reactive behavior when mixed with NH3 . Tertiary mixtures with
CO and CO2 , two other important constituents in interstellar ices, have been measured
(see Appendix 3.6) but here the differences are small compared with the observed binary
water-rich or CO-rich mixtures, and do not offer an alternative explanation.
3.4 Comparison between astronomical and laboratory
data
3.4.1 8–10 µm range
The FWHM and band positions of the NH3 and CH3 OH features measured in the laboratory and astronomical spectra are shown in Figs. 3.11 (for NH3 ) and 3.12 (for CH3 OH).
For the YSOs, the values obtained after removal of the silicate absorption (see §3.2) using the local continuum method are indicated by filled squares, whereas those obtained
from the template method are plotted with open squares. Note that the presence of significant amounts of CH3 OH may artificially lower the inferred NH3 ν2 width in CH3 OH
rich sources (indicated with * in Table 3.2) because of the contribution of the narrower ν7
CH3 -rock mode.
Regardless of the method used to subtract the continuum, or the type of source
(CH3 OH-rich/poor), we find that the observational band positions of the ν2 NH3 umbrella mode absorptions vary, within the errors, between 8.9 and 9.1 µm. This position
is not well reproduced by any of the investigated mixtures, but the positions measured
in water-rich ice mixtures are the closest, whereas the positions in pure NH3 or CO/CO2
rich ices are too far away to be representative of the astronomical positions. The derived
Spitzer FWHM values range between 0.23 and 0.32 µm (except for B1-b : 0.39 µm), when
using the local continuum method, not depending on whether the target is CH3 OH-rich or
-poor. For the template method, CH3 OH-rich sources generally tend to have a narrower
inferred FWHM, 0.3–0.5 µm, contrary to what would be expected if the NH3 mode is contaminated by the CH3 -rock feature. In any case, most of these widths are still narrower
than the laboratory FWHM values. To investigate further the effect of the continuum on
the positions and widths of the bands, we performed the following calculation to check
whether a continuum could be found that would yield NH3 and CH3 OH features with parameters within the laboratory measurements. To do that, we fitted the data between 8.25
and 10.4 µm with a function that is the sum of a 4th order polynomial and two Gaussians;
56
3.4 Comparison between astronomical and laboratory data
positions and widths of the Gaussians were constrained with limits taken from the laboratory data of binary water mixtures (8.9–8.95 µm for the NH3 position, 0.42–0.52 µm for
its width; 9.67–9.77 µm for the CH3 OH position, 0.2–0.3 µm for its width). We found
that the continuum derived in this way is different from those determined via the other
two methods. This result supports the fact that the difference between astronomical and
laboratory data could be attributed to the uncertainty in the continuum determination.
Taking the above considerations into account, Figs. 3.11 and 3.12 suggest that the
template method for subtraction of the 10 µm silicate absorption is more consistent with
the laboratory measurements, but both methods probably miss some weak NH3 absorption
features in the broad line wings where they blend with the continuum at the S /N of the
data. If so, the too small line widths inferred from the data (most probably due to the
uncertainty in the continuum determination) would mean that we have underestimated
NH3 abundances by a up to a factor of 2.
The observational band position and FWHM of the CH3 OH features derived with either the local continuum or template method are clustered around 9.7–9.75 µm, with the
exception of R CrA IRS 5 at 9.66 µm. Similarly the FWHM of the CH3 OH features
are all very similar between ∼0.22 and 0.32 µm, except for R CrA IRS 5 with 0.39 µm.
These values agree (with a few exceptions) with the values obtained from the laboratory
spectra. Note that the observed positions of the CH3 OH feature are all on the low side of
the laboratory range. Since the position of this feature shifts to higher wavelengths with
increasing water content, the observed low values could therefore indicate that CH3 OH
and H2 O are not well mixed and that there exists a separate CH3 OH-rich component, as
suggested in previous work [e.g. Pontoppidan et al. 2003a, Skinner et al. 1992]. Alternatively, the low values could be due to the presence of CO as indicated by the CH3 OH
feature shift to 9.70 µm in CH3 OH:CO=1:1. Both interpretations would be consistent
with the bulk of the CH3 OH formation coming from hydrogenation of a CO-rich layer,
rather than photochemistry in a water-rich matrix. However, the shift from the waterrich mixtures is small, and some water-rich fraction cannot be excluded with the current
spectral resolution.
3.4.2 The 3 and 6 µm ranges
Dartois & d’Hendecourt [2001] discussed the possibility of a 3.47 µm absorption band
which could be related to the formation of an ammonia hydrate in the ice mantles: they
found that if this band were mostly due to this hydrate, then ammonia abundances would
be at most 5% with respect to water ice. Considering the fact that our derived abundances
are larger than 10% in some sources, it is necessary to investigate the effect of such a
high abundance on the ammonia features in other spectral ranges. For this, depending
on the NH3 -to-CH3 OH abundance ratio observed in the Spitzer spectra, we scale one of
the following laboratory spectra to the 9 µm NH3 feature: H2 O:NH3 =9:1, H2 O:NH3 =4:1,
H2 O:CH3 OH:NH3 =10:1:1, H2 O:CH3 OH:NH3 =10:4:1. Figure 3.13 illustrates the comparison between the Spitzer and scaled laboratory spectra for the relevant wavelength
ranges for a couple of sources, while Figs. 3.14 and 3.15 (see Appendix 3.6) show the
57
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Figure 3.11 FWHM and band maximum position of the NH3 feature measured in the
laboratory mixtures at 15 K (“Lab.”, top panel) and in the Spitzer spectra (“Astro.”,
bottom panel). In the top panel, the filled star indicates pure NH3 , filled circles represent H2 O-rich mixtures and filled triangles are for NH3 :CH3 OH mixtures (an increasing symbol size indicative of increasing CH3 OH content). Other symbols are as follows: + for NH3 :H2 O=1:0.11, ▽ for NH3 :H2 O=1:1, ⋄ for NH3 :H2 O:CO=1:1:1, △ for
NH3 :H2 O:CO2 =1:1:1, for NH3 :CO:CO2 =1:1:1, × for NH3 :CH3 OH:H2 O=1:1:1. In the
bottom panel, open and filled squares indicate values obtained with the template and local
continuum method, respectively. The dash-dot polygons delimitate the parameter space
of FWHM and positions corresponding to H2 O-rich mixtures.
comparison for all sources where NH3 was detected. For further comparison, in Appendix 3.6 we also overplotted in Figs. 3.14 and 3.15 the following spectra: (i) the pure
H2 O ice spectrum derived from the H2 O column density quoted in Boogert et al. [2008]
(deep blue); and (ii) for sources with 3 µm data, the pure H2 O spectrum scaled to the
optical depth of the 3 µm feature of the mixed ice laboratory spectrum (purple-dotted).
The difference between this scaled pure water spectrum and the mixed ice spectrum gives
an indication of the contribution of ammonia features around 3.47 and 6.1 µm.
We then determined the contributions from the NH3 features to the integrated optical
depths of the 3 and 6 µm bands and to the optical depth of component C2, a feature at
58
3.4 Comparison between astronomical and laboratory data
Figure 3.12 Same as Fig. 3.11 but for CH3 OH. In the top panel, the filled star is for
pure CH3 OH, the filled square is representative of a CO-rich mixture. All other symbols (top and bottom panels) have the same meaning as in Fig. 3.11, except for the following in the top panel: + for NH3 :CH3 OH:H2 O=1:1:1, ▽ for CH3 OH:H2 O=1:1, ⋄ for
CH3 OH:CO=1:1, △ for CH3 OH:CO=9:1.
6.0-6.4 µm arising from a blend of several species, including NH3 , H2 O, CO2 , HCOO−
(see Paper I for more details). These contributions are reported in Tables 3.3 and 3.4.
Figures 3.14 and 3.15 (Appendix 3.6), and Tables 3.3 and 3.4 show that (i) the scaled
laboratory spectra generally do not overestimate the observed absorption features, and (ii)
for most sources, the presence of NH3 at the level we determine from the 9 µm feature
does not explain by itself the depth of the C2 component and of the red wing of the 3 µm
band. Hence, our inferred NH3 abundances up to 15% from the 9.7 µm data are not in
conflict with the lack of other NH3 features. The only exceptions are two sources (RNO
15 and EC 74), for which the scaled mixed ice spectrum exceeds the data in the 3 µm
range. In the case of RNO 15, the NH3 abundance could have been overestimated due to
the contribution of the CH3 OH CH3 -rock feature at ∼9 µm. For EC 74, this overestimate
and the presence of emission weakens the identification of the NH3 signature. In both
cases, the quoted NH3 abundances should be considered as upper limits.
59
NH3 contribution to the 3 and 6 µm bands for sources with a template
Source
IRAS 03235+3004
IRAS 03254+3050
IRAS 04108+2803
HH 300
IRAS 08242-5050
IRAS 08242-5050
2MASSJ17112317-272431
SVS 4-5
R CrA IRS 5
R
τH O,3.0
R 2
τ3.0
–
0.73
0.70
0.70
0.76
0.76
–
0.91
0.85
R
τmix,3.0
R
τ3.0
–
1.30
0.67
0.57
0.72
0.56
–
0.94
0.42
R 1785
τ
1562 H2 O
R 1785
τ
1562
0.50
0.56
0.58
0.50
0.50
0.50
0.69
0.42
0.63
R 1785
τ
1562 mix
R 1785
τ
1562
0.24
0.92
0.53
0.39
0.45
0.35
0.53
0.29
0.29
R
τNH3 ,6.16
R 1785
τ
1562 H2 O
τNH3 ,6.16
0.02
0.12
0.06
0.05
0.06
0.05
0.05
0.00
0.03
0.61
1.72
0.49
0.45
0.46
0.36
4.23
0.08
0.21
τC2
Note. — A dash indicates that the ratio was not calculated due to the high noise in the 3 µm spectrum.
RParameters are:
τH2 O,3.0 = integrated optical depth of pure water at 3 µm, determined from the column density of paper I and a band
−16
−1
RstrengthR of 2.0×10 cm .
τ3.0 , τmix,3.0 = integrated optical depth over the entire 3 µm region for, respectively, the considered source and the
corresponding laboratory mixture (selected from the NH3 feature at 9 µm).
R 1785
R 1785 R 1785
τ
, 1562 τ, 1562 τmix = integrated optical depth of, respectively, pure water, source spectrum, and laboratory
1562 H2 O
mixture,
between 1562 and 1785 cm−1 (5.6 to 6.4 µm).
R
τNH3 ,6.16 , τNH3 ,6.16 = integrated and peak optical depth of the 6.16 µm feature of ammonia obtained after subtraction of
a pure water spectrum scaled to the optical depth at 3 µm of the laboratory mixture.
τC2 = peak optical depth of the C2 component from paper I.
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
60
Table 3.3.
Table 3.4.
NH3 contribution to the 3 and 6 µm bands for sources with no associated
template
Source
τH O,3.0
R 2
τ3.0
0.80
–
–
0.78
0.94
0.65
0.95
0.90
–
–
R
τmix,3.0
R
τ3.0
1.97
–
–
0.88
0.94
0.95
2.34
0.35
–
–
R 1785
τ
1562 H2 O
R 1785
τ
1562
0.53
0.36
0.67
0.60
0.53
0.62
0.57
0.38
0.81
0.62
R 1785
τ
1562 mix
R 1785
τ
1562
1.23
0.44
0.43
0.56
0.45
0.75
1.18
0.10
0.19
0.55
R
τNH3 ,6.16
R 1785
τ
1562 H2 O
τNH3 ,6.16
0.16
0.05
0.03
0.04
0.04
0.05
0.09
0.00
0.00
0.06
0.45
0.60
0.57
0.83
0.53
1.64
0.76
0.01
0.08
0.34
τC2
Note. — A dash indicates that the ratio was not calculated due to the high noise in the 3 µm spectrum.
RParameters are:
τH2 O,3.0 = integrated optical depth of pure water at 3 µm, determined from the column density of paper I and a
of 2.0×10−16 cm−1 .
Rband strength
R
τ3.0 , τmix,3.0 = integrated optical depth over the entire 3 µm region for, respectively, the considered source and
the corresponding laboratory mixture (selected from the NH3 feature at 9 µm).
R 1785
R 1785 R 1785
τ
, 1562 τ, 1562 τmix = integrated optical depth of, respectively, pure water, source spectrum, and
1562 H2 O
laboratory
mixture, between 1562 and 1785 cm−1 (5.6 to 6.4 µm).
R
τNH3 ,6.16 , τNH3 ,6.16 = integrated and peak optical depth of the 6.16 µm feature of ammonia obtained after
subtraction of a pure water spectrum scaled to the optical depth at 3 µm of the laboratory mixture.
τC2 = peak optical depth of the C2 component from paper I.
61
3.4 Comparison between astronomical and laboratory data
RNO 15
IRAS 03271+3013
B1-a
L1489 IRS
RNO 91
IRAS 17081-2721
EC 74
EC 92
CrA IRS7 B
L1014 IRS
R
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Overall, our reported NH3 abundances are up to a factor of three larger than the upper
limits derived by Dartois & d’Hendecourt [2001]. Firstly, let’s recall that the conclusions
in their study and in ours are drawn from the analysis of different samples. Secondly, Dartois & d’Hendecourt made an assumption that does not apply to our sample: indeed, they
considered a grain size distribution including also scattering from larger grains, producing
an enhanced 3 µm wing, whereas the results presented here can be taken as representative
of NH3 absorption from small grains. It is beyond the scope of this paper to investigate
the effects of grain size distribution and scattering in as much detail as did Dartois &
d’Hendecourt [2001].
Figure 3.13 Comparison of astronomical data (VLT or Keck measurements at short wavelengths, IRS Spitzer observations elsewhere) and laboratory spectra in selected wavelength ranges: 2.0–4.5 µm (left panels), 5.2–7.5 µm (middle panels) and 8.2–10.2 µm
(right panels, silicate absorption subtracted via the template method). Error bars are
indicated in the bottom-right corner. Overlaid in red and green are laboratory spectra
corresponding to H2 O:CH3 OH:NH3 =10:4:1 and H2 O:NH3 =9:1, respectively, scaled to
the 9 µm NH3 umbrella mode. The dark blue line represents the pure water laboratory
spectrum scaled to the water column density taken in paper I. The dotted purple line
corresponds to a pure water spectrum scaled to the 3 µm water feature of the mixed ice
spectrum, showing the contribution of NH3 features around 3.47 and 6.1 µm. Finally, the
red dashed line in the right panel of SVS 4-5 represents a H2 O:CH3 OH=9:1 laboratory
spectrum scaled to the 9.7 µm CH3 OH CO-stretch mode: this gives an indication of the
contribution of the 9 µm CH3 OH CH3 -rock mode to the total 9 µm feature. The laboratory
spectra are recorded at 15 K unless indicated differently.
62
3.5 Conclusion
3.4.3 Nitrogen ice inventory
The confirmation of the presence of relatively large amounts of solid NH3 , up to 15%, in
interstellar ices solves a long-standing problem. Indeed, the detection of solid NH3 has
remained elusive and/or controversial, despite a number of clues suggesting its presence:
• High cosmic abundance of atomic nitrogen : NN /NH = 7.76 × 10−5 [Savage &
Sembach 1996], only a factor of a few below those of oxygen and carbon. Here NH
indicates the total number of hydrogen nuclei, NH =N(H)+2N(H2 ).
• High abundances of gaseous NH3 of NNH3 /NH2 ∼ 10−6 − 10−5 in the Orion-KL
nebula [Barrett et al. 1977, Genzel et al. 1982] and in other hot cores such as
G9.62+0.19, G29.96−0.02, G31.41+0.31 [Cesaroni et al. 1994], and G10.47+0.03
[Cesaroni et al. 1994, Osorio et al. 2009].
• Identification of substantial amounts of OCN− [e.g. van Broekhuizen et al. 2004,
2005] and NH+4 in ices [e.g. Schutte & Khanna 2003, Boogert et al. 2008]: considering that these ions form via reactions involving NH3 , the non-detection of solid
NH3 would be puzzling.
Our results can be used to draw up a possible nitrogen budget. Assuming NH2 O /NH ∼
5×10−5 [Pontoppidan et al. 2004, Boogert et al. 2004], and average abundances w.r.t. H2 O
of 5.5% for NH3 (see §3.2.3), 7% for NH+4 (from Table 3 of Paper I), and 0.6% for OCN−
[van Broekhuizen et al. 2005], then the NH3 , NH+4 and OCN− abundances with respect
to total H are 2.8, 3.5, and 0.3 ×10−6 respectively. This corresponds to, respectively, 3.4,
4.4 and 0.4% of the atomic nitrogen cosmic abundance so that, in total, about 10% of the
cosmically available nitrogen would be locked up in ices, leaving solid and gaseous N2 , N
and HCN as other substantial nitrogen carriers. The main uncertainty in this determination
is the adopted H2 O ice abundance with respect to total H; in several sources this may well
be a factor of 2 larger, leading to about 20% of the nitrogen accounted for in ices.
3.5 Conclusion
We have analyzed in detail the 8-10 µm range of the spectra of 41 low-mass YSOs obtained with Spitzer and presented in Boogert et al. [2008]. The sources are categorized
into three types: straight, curved and rising 8 µm silicate wings, and for each category
template sources with little or no absorption from ices around 9–10 µm have been determined. This has led to two ways of subtracting the contribution from the 10 µm silicate
absorption: first, by determining a local continuum, and second, by scaling the templates
to the optical depth at 9.7 µm. The two methods give consistent band positions of the NH3
features, but the resulting widths can be up to a factor of two larger using the template
continuum method. Taking into account the uncertainty in continuum determination, NH3
ice is most likely detected in 24 of the 41 sources with abundances of ∼2 to 15% w.r.t.
H2 O, with an average abundance of 5.5±2.0%. These abundances have estimated uncertainties up to a factor of two and are not inconsistent with other features in the 3 and 6 µm
63
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
ranges. CH3 OH is often detected as well, but the NH3 /CH3 OH abundance ratio changes
strongly from source to source. Our inferred CH3 OH column densities are consistent with
the values derived in paper I.
Targeted laboratory experiments have been carried out to characterize the NH3 and
CH3 OH profiles (position, FWHM, integrated absorbance). Comparison with the observational data shows reasonable agreement (within ∼1%) for the position of the NH3
feature in H2 O-rich ices, but the observed widths are systematically smaller than the laboratory ones for nearly all sources. The silicate template continuum method gives widths
that come closest to the laboratory values. This difference in width (i.e. widths derived
from astronomical spectra smaller than those in the laboratory spectra) suggests that the
NH3 abundances determined here may be on the low side.
The CH3 OH profile is most consistent with a significant fraction of the CH3 OH in a
relatively pure or CO-rich phase, consistent with its formation by the hydrogenation of
CO ice. In contrast, the most likely formation route of NH3 ice remains hydrogenation of
atomic N together with water ice formation in a relatively low density molecular phase.
Finally, the nitrogen budget indicates that up to 10 to 20% of nitrogen is locked up in
known ices.
3.6 Appendix
Parameters of Gaussian fits
64
Table 3.5.
(a) Parameters of Gaussian fits to the NH3 feature.
Source
λ (µm)
IRAS 03235+3004
L1455 IRS3
IRAS 03254+3050
B1-b∗
IRAS 04108+2803
HH 300
IRAS 08242-5050
IRAS 15398-3359
B59 YSO5
2MASSJ17112317-272431
SVS 4-5∗
R CrA IRS 5
8.93±0.02
8.99±0.03
9.04±0.01
9.05±0.03
8.99±0.02
9.01±0.02
9.02±0.01
8.96±0.01
8.95±0.01
8.99±0.01
9.00±0.01
9.05±0.02
NH3 , local
FWHM (µm)
0.28±0.03
0.24±0.07
0.25±0.03
0.39±0.06
0.25±0.04
0.23±0.05
0.31±0.03
0.29±0.03
0.27±0.03
0.30±0.02
0.26±0.03
0.21±0.04
τpeak
λ (µm)
0.23±0.02
0.02±0.01
0.10±0.01
0.25±0.02
0.05±0.01
0.04±0.01
0.15±0.01
0.30±0.02
0.18±0.02
0.43±0.02
0.16±0.02
0.04±0.01
8.93±0.01
9.02±0.02
8.99±0.01
9.07±0.03
9.05±0.03
9.06±0.02
9.05±0.01
8.98±0.01
8.89±0.02
9.02±0.02
9.01±0.01
9.00±0.03
NH3 , template
FWHM (µm)
0.29±0.03
0.38±0.05
0.38±0.03
0.40±0.06
0.47±0.06
0.45±0.06
0.30±0.03
0.33±0.03
0.34±0.04
0.50±0.05
0.30±0.03
0.36±0.06
τpeak
0.30±0.02
0.04±0.01
0.12±0.01
0.31±0.03
0.04±0.01
0.05±0.01
0.15±0.01
0.41±0.02
0.18±0.02
0.41±0.04
0.26±0.02
0.04±0.01
3.6 Appendix
65
Source
λ (µm)
RNO 15
IRAS 03271+3013
B1-a
L1489 IRS
IRAS 13546-3941
RNO 91
IRAS 17081-2721
EC 74
EC 82
EC 90
EC 92∗
CK4
CrA IRS7 B∗
L1014 IRS
9.05±0.02
8.96±0.02
8.98±0.02
9.02±0.01
8.99±0.02
8.98±0.01
8.97±0.02
9.01±0.02
8.94±0.01
8.95±0.02
8.99±0.02
8.99±0.02
9.04±0.01
9.03±0.02
NH3 , local
FWHM (µm)
0.20±0.04
0.25±0.04
0.25±0.04
0.27±0.03
0.27±0.03
0.26±0.03
0.24±0.04
0.22±0.05
0.29±0.03
0.18±0.05
0.25±0.05
0.32±0.04
0.26±0.03
0.25±0.05
(a) Cont’d
τpeak
λ (µm)
0.04±0.01
0.20±0.02
0.14±0.02
0.09±0.01
0.03±0.00
0.08±0.01
0.04±0.00
0.05±0.01
0.04±0.00
0.04±0.01
0.03±0.00
0.03±0.00
0.15±0.01
0.15±0.02
–
–
–
–
–
–
–
–
–
–
–
–
–
–
NH3 , template
FWHM (µm)
–
–
–
–
–
–
–
–
–
–
–
–
–
–
τpeak
–
–
–
–
–
–
–
–
–
–
–
–
–
–
Note. — Uncertainties are statistical errors from the Gaussian fits.
∗
Sources with τ9.7µm > 2 × τ9.0µm , for which the contribution from the CH3 OH CH3 -rock mode is significant.
Since the latter and the NH3 umbrella mode were difficult to disentangle, a single fit was performed (the reported
parameters) and the integrated optical depth of the ammonia feature was then obtained from the total integrated
optical depth at 9 µmby subtracting the estimated contribution of the CH3 OH CH3 -rock mode (see §3.2.2).
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
66
Table 3.5.
Table 3.5. (b) Parameters of Gaussian fits to the CH3 OH C-O stretch mode (after
subtraction of the continuum with the local and/or template method), and CH3 OH
column densities (or 3-σ upper limits).
Source
λ
(µm)
IRAS 03235+3004
L1455 IRS3
IRAS 03254+3050
B1-b
IRAS 04108+2803
HH 300
IRAS 08242-5050
IRAS 15398-3359
B59 YSO5
2MASSJ17112317-272431
SVS 4-5
R CrA IRS 5
9.74±0.02
9.78±0.01
···
9.71±0.01
···
···
9.70±0.01
9.73±0.01
···
9.75±0.02
9.74±0.01
9.66±0.01
Local continuum
FWHM
τpeak
(µm)
0.26±0.03
0.14±0.03
···
0.30±0.03
···
···
0.27±0.03
0.28±0.03
···
0.23±0.04
0.28±0.03
0.39±0.03
0.35±0.04
0.03±0.01
···
1.19±0.11
···
···
0.25±0.02
0.77±0.06
···
0.13±0.02
0.77±0.06
0.07±0.00
X
(% H2 O)
λ
(µm)
4.40±1.04
3.67±1.80
< 5.4
14.15±3.16
< 2.7
< 4.7
6.12±1.01
10.26±3.02
< 1.2
1.03±0.22
26.38±6.17
5.68±0.60
9.74±0.02
9.78±0.02
···
9.71±0.01
9.75±0.00
9.74±0.00
9.70±0.01
9.73±0.01
···
···
9.74±0.01
9.66±0.02
Template continuum
FWHM
τpeak
(µm)
0.25±0.04
0.26±0.04
···
0.28±0.03
0.06±0.04
0.19±0.12
0.29±0.03
0.30±0.03
···
···
0.31±0.02
0.39±0.04
0.31±0.04
0.04±0.01
···
1.21±0.11
0.04±0.03
0.01±0.01
0.24±0.02
0.75±0.06
···
···
0.83±0.06
0.07±0.00
X
(% H2 O)
Paper I
X
(% H2 O)
3.84±0.99
7.71±3.46
< 5.4
13.75±3.12
0.58±0.62
0.78±0.52
6.39±1.09
10.69±3.14
< 1.2
< 2.0
31.50±7.12
5.51±0.72
4.20±1.20
<12.5
< 4.6
11.20±0.70
< 3.5
< 6.7
5.50±0.30
10.30±0.80
< 1.3
< 3.2
25.20±3.50
6.60±1.60
3.6 Appendix
67
Source
λ
(µm)
RNO 15
IRAS 03271+3013
B1-a
L1489 IRS
IRAS 13546-3941
RNO 91
IRAS 17081-2721
EC 74
EC 82
EC 90
EC 92
CK4
CrA IRS7 B
L1014 IRS
9.65±0.03
···
···
9.78±0.02
···
9.77±0.01
···
···
···
9.70±0.01
9.73±0.01
···
9.70±0.01
9.69±0.03
Local continuum
FWHM
τpeak
(µm)
0.44±0.07
···
···
0.10±0.03
···
0.11±0.03
···
···
···
0.32±0.03
0.30±0.02
···
0.33±0.02
0.38±0.08
0.02±0.00
···
···
0.03±0.01
···
0.05±0.01
···
···
···
0.05±0.00
0.09±0.01
···
0.36±0.02
0.10±0.01
(b) Cont’d.
X
(% H2 O)
λ
(µm)
11.13±2.16
< 4.3
< 2.4
0.44±0.22
< 2.0
0.87±0.32
< 6.6
<13.5
<24.6
6.91±0.99
11.16±1.46
···
7.74±1.56
3.61±0.99
···
···
···
···
···
···
···
···
···
···
···
···
···
···
Template continuum
FWHM
τpeak
X
(µm)
(% H2 O)
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
Paper I
X
(% H2 O)
< 5.0
< 5.6
< 1.9
4.90±1.50
< 3.9
< 5.6
3.30±0.80
< 9.3
<14.2
6.80±1.60
11.70±3.50
···
6.80±0.30
3.10±0.80
Note. — This table shows that CH3 OH column densities obtained in this paper are consistent with those in Paper I, which are our recommended
values.
Note. — Uncertainties are statistical errors from the Gaussian fits.
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
68
Table 3.5.
3.6 Appendix
Additional laboratory data
The table presented here (Table 3.6) gives an overview of all the ice mixtures measured
in the laboratory. The table contains only spectra taken at a sample temperature of 15 K.
The focus in the table is on the band position of the NH3 ν2 umbrella mode, the CH3 OH
ν4 C-O stretch mode and the CH3 OH ν7 CH3 rock mode. Full width at half maximum
(FWHM) and the position of the maximum of the absorption profile of these modes are
indicated in both cm−1 and µm for each mixture. For the NH3 ν2 umbrella mode, the band
strength relative to that of the pure NH3 mixture is also indicated.
69
NH3
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
a
Ice mixture
CH3 OH
H2 O
0
0
0
0
0
0
0
0
0
4
4
4
4
2
2
2
2
1
1
1
1
0.5
0.5
0.5
0.5
0.25
0.25
0.25
0.25
0
0.11
1
9
10
1
10
1
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
CO
0
0
0
0
1
1
0
0
1
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
Molecule
Peak position
cm−1
µm
FWHM
cm−1
µm
Arel.
NH3
NH3
NH3
NH3
NH3
NH3
NH3
NH3
NH3
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
1070
1076
1100
1118
1124
1094
1122
1098
1062
1129a
1029
1128
2823
1111a
1029
1132
2820
1086
1029
1135
2817
1080
1030
1128a
2813
1078
1030
–a
2808a
66
70
77
62
53
75
57
82
66
108a
30
35
28
115a
29
35
26
137
26
44
26
118
22
35a
27
98
16
–a
17a
1
1
1
0.7
0.7
1
0.8
0.9
0.8
0.4a
–
–
–
0.6a
–
–
–
0.8
–
–
–
0.8
–
–
–
0.9
–
–
–a
CO2
0
0
0
0
0
0
2
1
1
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
Band is weak and spectral overlap prohibits accurate fitting.
9.341
9.291
9.091
8.947
8.897
9.144
8.916
9.108
9.414
8.856a
9.722
9.707
3.543
8.994a
9.720
8.833
3.546
9.209
9.716
8.813
3.550
9.258
9.711
8.865a
3.555
9.278
9.707
–a
3.561a
0.577
0.605
0.637
0.496
0.420
0.627
0.453
0.681
0.586
0.849
0.283
0.275
0.035
0.934
0.274
0.273
0.033
1.166
0.246
0.342
0.033
1.015
0.207
0.275
0.034
0.845
0.151
–a
0.022
Mode
ν2
ν2
ν2
ν2
ν2
ν2
ν2
ν2
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
umbrella
umbrella
umbrella
umbrella
umbrella
umbrella
umbrella
umbrella
umbrella
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
70
Table 3.6. Ice composition, band maximum position (“peak position”), FWHM and
band strength relative to the pure ice (Arel. ), listed for a set of ice mixtures under
investigation.
Table 3.6.
NH3
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
1
0
0
0
0
0
0
0
0
0
1
1
1
1
0.25
0.25
0.25
0.25
1
1
1
1
4
4
4
4
1
1
1
1
1
1
1
1
1
1
1
1
1
10
10
10
10
10
10
10
10
10
10
10
10
0
0
0
1
1
1
9
9
9
CO
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
Molecule
Peak position
cm−1
µm
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
NH3
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
CH3 OH
1116a
1026
1125a
2824
1119
1017
–a
2829a
1123
1022
–a
2830
1130
1023
1124
2830
1028
1125
2828
1025
1124
2828
1020
1126
2828
CO2
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
0
Band is weak and spectral overlap prohibits accurate fitting.
8.961
9.745
8.888a
3.541
8.937
9.833
–a
3.534a
8.903
9.784
–a
3.533
8.848
9.777
8.896
3.534
9.729
8.888
3.536
9.755
8.897
3.536
9.801
8.883
3.536
FWHM
cm−1
µm
95
29
32a
26
59
22
–a
30a
61
24
–a
15
62
30
23
14
28
34
33
33
40
23
23
13
23
0.764
0.276
0.253
0.033
0.472
0.213
–a
0.037
0.484
0.230
–a
0.019
0.489
0.288
0.183
0.017
0.265
0.269
0.041
0.314
0.317
0.029
0.221
0.103
0.029
Arel.
0.7
–
–
–
1
–
–
–
1
–
–
–
–
–
–
–
1
1
1
–
–
–
–
–
–
Mode
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν2
ν4
ν7
ν2
ν4
ν7
ν2
ν4
ν7
ν2
ν4
ν7
ν2
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
umbrella
C-O stretch
CH3 rock
C-H stretch
C-O stretch
CH3 rock
C-H stretch
C-O stretch
CH3 rock
C-H stretch
C-O stretch
CH3 rock
C-H stretch
71
3.6 Appendix
a
Ice mixture
CH3 OH
H2 O
(cont’d)
NH3
0
0
0
0
0
0
0
0
0
0
0
0
a
Ice mixture
CH3 OH
H2 O
1
1
1
1
1
1
1
1
9
9
9
9
0
0
0
0
0
0
0
0
0
0
0
0
CO
9
9
9
9
1
1
1
1
1
1
1
1
Molecule
Peak position
cm−1
µm
CH3 OH
CH3 OH
CH3 OH
CO
CH3 OH
CH3 OH
CH3 OH
CO
CH3 OH
CH3 OH
CH3 OH
CO
1034
1119
2831
2138
1029
1124
2830
2136
1028
1125
2824
2135
CO2
0
0
0
0
0
0
0
0
0
0
0
0
(cont’d)
Band is weak and spectral overlap prohibits accurate fitting.
9.675
8.938
3.532
4.677
9.720
8.898
3.534
4.682
9.730
8.890
3.541
4.685
FWHM
cm−1
µm
25
30
–
7
30
32
–
9
28
32
–
9
0.229
0.242
–
0.014
0.286
0.258
–
0.020
0.261
0.255
–
0.021
Arel.
–
–
–
–
–
–
–
–
–
–
–
–
Mode
ν4
ν7
ν2
ν1
ν4
ν7
ν2
ν1
ν4
ν7
ν2
ν1
C-O stretch
CH3 rock
C-H stretch
C-O stretch
C-O stretch
CH3 rock
C-H stretch
C-O stretch
C-O stretch
CH3 rock
C-H stretch
C-O stretch
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
72
Table 3.6.
3.6 Appendix
Comparison between astronomical and laboratory data
Comparison between astronomical and laboratory data for sources whose silicate absorption feature was fitted with a template. For a given source (displayed in either the left or
right column of the figure), the middle and right panels show 5.2–7.5 and 8.2–10.2 µm
regions from IRS Spitzer spectra overlaid with laboratory spectra, scaled to the 9 µm NH3
umbrella mode. Error bars for the Spitzer spectra are indicated in the bottom-right corner. The dark blue line represents the pure water laboratory spectrum scaled to the water
column density taken in paper I. Other colors are representative of laboratory spectra obtained for the following mixtures: H2 O:NH3 =9:1 (green), H2 O:CH3 OH:NH3 =10:0.25:1
(orange), H2 O:CH3 OH:NH3 =10:1:1 (cyan), and H2 O:CH3 OH:NH3 =10:4:1 (red). When
available (see Boogert et al. 2008), VLT or Keck data (2.0–4.5 µm, left panel) are also
plotted. In this case, we overplotted (dotted purple line) a pure water spectrum scaled to
the 3 µm water feature of the mixed ice spectrum. Whenever present, a red dashed line
in the right panel of a given source represents a H2 O:CH3 OH=9:1 laboratory spectrum
scaled to the 9.7 µm CH3 OH CO-stretch mode: this gives an indication of the contribution
of the 9 µm CH3 OH CH3 -rock mode to the total 9 µm feature. The laboratory spectra are
recorded at 15 K unless indicated differently. Please refer to the on-line version for the
color coding.
73
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Figure 3.14 (a) Caption as described on the previous page
74
3.6 Appendix
Figure 3.15 same as Fig. 3.14.
75
3 The c2d spectroscopic survey of ices. IV NH3 and CH3 OH
Figure 3.15 (b) As for Fig. 3.14 but for sources with no associated template, i.e. with the
10 µm silicate feature subtracted via the local continuum method. Additionally, yellow
represents H2 O:NH3 =4:1 (H. Fraser, priv. comm.).
76
3.6 Appendix
Figure 3.15 (b) As (a) but for sources with no associated template, i.e. with the 10 µm silicate feature subtracted via the local continuum method. Additionally, yellow represents
H2 O:NH3 =4:1 (H. Fraser, priv. comm.).
77
CHAPTER 4
IR spectroscopy of VUV irradiated PAH
containing interstellar ices1
Polycyclic aromatic hydrocarbons (PAHs) are known to be abundantly present in photondominated regions (PDRs), as evidenced by their ubiquitous mid-IR emission bands. Towards dense clouds, however, their IR emission bands are strongly suppressed. It is here
where molecules are known to reside on very cold grains (T ≤30 K) in the form of interstellar ices. Therefore, it is likely that non-volatile species, such as PAHs, also freeze out
on grains. Such icy grains act as catalytic sites and, upon vacuum ultraviolet (VUV) irradiation, chemical reactions are initiated. These reactions and the resulting photoproducts
are investigated in the study presented here for PAH containing water ices. The aim of
this work is to monitor vacuum ultraviolet induced chemical reactions of PAHs in cosmic
ice through their IR signatures, to characterize the families of species formed in these
reactions, and to apply the results to astronomical observations. Mid-infrared Fourier
transform absorption spectroscopic measurements ranging from 6500 to 450 cm−1 are
performed on freshly deposited and vacuum ultraviolet processed PAH containing cosmic H2 O ices at low temperatures. The mid-IR spectroscopy of anthracene, pyrene and
benzo[ghi]perylene containing H2 O ice is reported. Band strengths of the neutral PAH
modes in H2 O ice are derived. Additionally, spectra of vacuum ultraviolet processed PAH
containing H2 O ices are presented. These spectra are compared to spectra measured in
VUV processed PAH:argon matrix isolation studies. It is concluded that the parent PAH
species is ionized in H2 O ice and that other photoproducts, mainly more complex PAH
derivatives, also form. The importance of PAHs and their PAH:H2 O photoproducts in astronomical mid-infrared spectroscopic studies, in particular in the 5–8 µm region, is discussed. As a test-case, the VUV photolyzed PAH:H2 O laboratory spectra are compared
to a high resolution ISO-SWS spectrum of the high-mass embedded protostar W33A and
to a Spitzer spectrum of the low-mass Young Stellar Object (YSO) RNO 91. For these
objects, an upper limit of 2–3% with respect to H2 O ice is derived for the contribution of
PAHs and PAH:H2 O photoproducts to the absorbance in the 5–8 µm region towards these
objects.
1 Based on: J. Bouwman, A. L. Mattioda, H. Linnartz, and L. J. Allamandola, Astronomy and Astrophysics,
submitted (2010)
79
4 IR spectroscopy of PAH containing ices
4.1 Introduction
Polycyclic aromatic hydrocarbons (PAHs) are known to be abundantly present in photondominated regions (PDRs) [Peeters et al. 2004a, van Dishoeck 2004, Tielens 2008]. The
evidence for the ubiquity of astronomical PAHs is the widespread, well-known family of
prominent emission bands at 3.28, 6.2, 7.6, 8.6, and 11.2 µm (3050, 1610, 1300, 1160,
and 890 cm−1 ) associated with many, if not most, galactic and extragalactic objects [Smith
et al. 2007, Draine & Li 2007]. These bands dominate the mid-IR emission spectrum because of an intrinsically high efficiency of the fluorescent process and are most easily
detected in regions where individual gas-phase PAH molecules (both neutrals and ions)
become highly vibrationally excited by the ambient UV-VIS-NIR radiation field [Mattioda et al. 2005a, Li & Draine 2002]. They then energetically relax by emission of IR
photons at frequencies corresponding to fundamental vibrational modes, resulting in these
well known emission spectra.
PAHs and related aromatic materials are expected to be present both in optically thin,
diffuse regions of the ISM and in dense environments. In dense regions, however, the
highly efficient PAH fluorescence is found to be quenched. There are two reasons for
this. First, the radiation which pumps the emission tapers off with extinction into dense
regions, and second, in cold molecular clouds PAHs can serve as nucleation sites on
which other species condense. In this way, neutral and/or charged PAHs can agglomerate
to form (charged) PAH clusters, or very small grains (VSGs) [e.g., Allamandola et al.
1989, Rapacioli et al. 2006]. The VSGs can, subsequently, freeze out on grains or serve
as nucleation sites for small molecules forming ice covered VSGs. Individual PAHs can
also efficiently condense onto dust grains as ‘guest molecules’ in icy grain mantles, much
as is the case for most other smaller interstellar molecules [e.g., Sandford & Allamandola
1993]. Vibrational energy of a PAH molecule which is part of a larger dust particle, either
as a nucleation center or guest in a water-rich ice, efficiently dissipates into the phonon
modes of the solid material on a time-scale orders of magnitude shorter than required
to emit an IR photon [Allamandola et al. 1985, 1989]. Consequently, in dark, dense
regions, PAHs and PAH derivatives are expected to give rise to IR absorption bands, not
to emission features.
There are several lines of evidence that support the presence of PAHs in dense molecular clouds. Aromatics in primitive meteorites and interplanetary dust particles contain
deuterium enrichments that are best explained by an interstellar cloud heritage [e.g., Sandford 2002,and references therein]. In addition, very weak absorption features attributed
to aromatic hydrocarbons have been observed in the IR absorption spectra of objects embedded in dense clouds. These include a band near 3.3 µm (3030 cm−1 ) [Smith et al.
1989, Sellgren et al. 1995, Brooke et al. 1999, Chiar et al. 2000], and bands near 6.2 µm
(1600 cm−1 ) [Chiar et al. 2000] and 11.2 µm (890 cm−1 ) [Bregman et al. 2000]. These
very weak features are severely blended with much stronger H2 O ice bands, consistent
with the number of PAH molecules relative to the number of H2 O molecules along these
lines of sight on the order of a few percent. So far, it has proven difficult to unambiguously
interpret these absorption features in spite of the fact that there is a growing database of
theoretically calculated and laboratory measured IR absorption spectra of both neutral
80
4.2 Experimental technique
and ionized PAHs in inert matrices [e.g., Szczepanski & Vala 1993, Szczepanski et al.
1993a,b, 1995a,b, Hudgins et al. 1994, Hudgins & Allamandola 1995a, 1997, Langhoff
1996, Mattioda et al. 2005b, Bauschlicher et al. 2009, 2010,and references therein]. Unfortunately, these spectra cannot be used directly to compare with PAHs in H2 O-rich
ices, as rare gas matrix spectra will be different. Intermolecular interactions perturb the
molecular vibrational energy levels, influencing IR band positions, widths, profiles, and
intrinsic strengths. Consequently, it has not yet been possible to properly evaluate astronomical solid state PAH features, mainly because the corresponding laboratory data of
realistic ice analogs are lacking.
Therefore, in the Astrochemistry Laboratory at NASA Ames Research Center a program to measure the IR spectra of PAHs in water ices was started. Earlier work focused
on the IR band positions, band widths, and relative band strengths of neutral PAHs [Sandford et al. 2004, Bernstein et al. 2005a,b]. More recently, an exploratory study of the
effects of vacuum ultraviolet (VUV) photolysis on several PAH:H2 O ice mixtures was
carried out [Bernstein et al. 2007]. Unfortunately, at concentrations that are most appropriate for dense clouds, PAH bands are swamped in the mid-IR by overlapping H2 O ice
bands and it has proven difficult to put these IR-only data on a solid quantitative footing.
This situation has changed thanks to the development of a new apparatus at the Sackler
Laboratory for Astrophysics at Leiden University which allows one to track the in situ,
VUV photochemistry of low concentration PAH:H2 O ices using optical (i.e. electronic)
spectroscopy (Chapter 5, 6, and 7). This approach makes it possible to simultaneously
follow the VUV driven kinetic behavior of the neutral parent PAH and photoproducts on
millisecond timescales with a concentration precision on the order of a few percent. Using this approach, it was possible characterized the VUV photochemistry of four PAHs
anthracene, pyrene, benzo[ghi]perylene, and coronene in water ice at various concentrations and ice temperatures. This is described in detail in Chapter 7. The mid-IR study
presented here is partially based on the quantitative results derived in the optical work.
This paper is laid out as follows. After the experimental technique is described in §7.2,
the results are presented in §4.3 and §4.4. These include band profiles and band strengths
for neutral anthracene, pyrene and benzo[ghi]perylene in H2 O ice, the IR spectroscopic
properties of their VUV induced photoproducts, and the visualization of photochemical
processes at play during extended photolysis. In §4.6 we extend our findings to the general IR properties of PAHs in ices and use these data to interpret observations of ices
in dense clouds towards the high-mass protostar W33A and the low-mass young stellar
object RNO 91. The conclusions are summarized in §4.7.
4.2 Experimental technique
The techniques employed in this study have been described in detail previously [Hudgins
et al. 1994] and the relevant details are summarized briefly. The ices are prepared by vapor
co-deposition of the PAH of interest with water vapor onto a 15 K CsI window which is
suspended in a high vacuum chamber (P≤ 10−8 Torr). The PAHs anthracene (Ant, C14 H10 ,
Aldrich, 99%) and pyrene (Py, C16 H10 , Aldrich, 99%) are used without further purifica81
4 IR spectroscopy of PAH containing ices
Table 4.1 Ice mixture (PAH:X, where X indicates the ice matrix species), PAH vaporization temperatures, resulting concentrations (PAH:Z, where Z indicates the relative amount
of the matrix species) for anthracene, pyrene, and benzo[ghi]perylene containing ices, and
the ice temperature during photolysis for the ice mixtures under investigation.
Ice (PAH:X)
Ant:H2 O
Py:H2 O
Py:CO
Bghi P:H2 O
Tdep (◦ C)
32
42
53
71
51
41
44
50
51
50
143
156
152
Conc. (PAH:Z)
1:450
1:172
1:60
1:11
1:100
1:200
1:90
1:65
1:70
1:30
1:160
1:60
1:110
Tice (K)
15
15
15
15
125
15
15
15
125
15
15
15
125
tion and vaporized from heated pyrex tubes. The PAH benzo[ghi]perylene (Bghi P, C22 H12 ,
Aldrich, 98%) is kept at a temperature of 180◦ C for 20 minutes with a cold shield blocking the deposition onto the sample window to remove most of the contaminants and is
subsequently deposited in a manner similar to that for Ant and Py. Simultaneously, water
vapor — milli-Q grade, further purified by three freeze-pump-thaw cycles — is admitted
through an adjacent deposition tube. To prepare PAH:H2 O ices with different concentrations, the PAH deposition temperature was varied from one experiment to the other, while
the water flow was kept constant. Mid-infrared spectroscopy of the PAHs Ant, Py, and
Bghi P in water ice and for Py:CO ice as a control experiment is performed for a range of
concentrations. The PAH deposition temperature, the resulting PAH concentration, and
the ice temperature during photolysis are summarized in Table 4.1.
VUV photolysis of the sample ices is accomplished with the combined 121.6 nm
Lyman-α (10.6 eV) and 160 nm (7.8 eV) molecular hydrogen emission bands from a
microwave powered discharge in a flowing H2 gas at a dynamic pressure of 150 mTorr.
The VUV radiation from the lamp enters the sample chamber through a MgF2 window.
The UV photon flux of the lamp is ∼1015 photons cm−2 s−1 at the sample surface.
Spectra from 6500 to 450 cm−1 are measured with a Biorad Excalibur FTS 4000 FTIR
spectrometer equipped with a KBr beamsplitter and a liquid N2 -cooled MCT detector.
Spectra are taken in optical depth (τ= ln I/I0 ), with the background spectrum (I0 ) taken
on the cold sample window before sample deposition, and the spectra (I) taken after
deposition and VUV processing of the sample. Each spectrum represents a co-addition of
512 spectra at a resolution of 0.5 cm−1 . This level of resolution is necessary to distinguish
photoproduct bands that are close to the position of the neutral band. The number of scans
was chosen to optimize both the signal-to-noise ratio as well as the time requirements for
82
4.2 Experimental technique
each experiment. Spectra are taken at 15 K of the freshly deposited sample and after 5,
10, 15, 30, 60, 120, and 180 minutes of VUV photolysis. Similar measurements are also
performed on PAH:H2 O ice samples at 125 K to check for differences in photochemistry
between low and high temperature photolysis experiments.
We focus on the bands in the 1650–1000 cm−1 region because this is where the PAH
mid-IR bands suffer the least from overlap with the strong H2 O ice features (Fig. 4.1).
Hereafter we refer to this as the Region Of Interest (ROI). The absolute band strengths for
the neutral PAH modes in H2 O ice are derived as follows. The H2 O bending and librational overtone modes are subtracted from the raw spectrum by fitting a spline function
through a set of points where no PAH absorption occurs. Subsequently, the total theoretically calculated absolute intensity in the ROI is proportionally divided over the measured
PAH modes in this frequency range, via:
 R νi,2
M
τ dν

X
νi,1 i,ν
exp 
thy
,
(4.1)
Ai = 
A j  PL R νi,2
i=0 ν τi,ν dν
j=0
i,1
exp
Ai
where
is the experimentally measured band strength of PAH mode i in H2 O ice in
cm molecule−1 , Athy
j is the theoretically calculated absolute intensity of vibrational mode
j in the ROI in cm molecule−1 , M is the number of theoretically calculated modes in the
ROI, τi,ν is the optical depth of mode i in H2 O ice at frequency ν (cm−1 ), L is the number of measured modes in the ROI, and νi,1 and νi,2 are the lower and upper integration
boundaries in cm−1 , respectively, for absorption feature i. This method takes advantage
of the fact that, although there may be band-to-band variations in the accuracy of the
calculated intensity for one band, the total theoretically calculated intensity is generally
accurate to within 10–20%. Here we assume that the matrix material does not substantially influence the total integrated IR band intensity and, as will be shown later, this is
only an approximation.
The PAH in H2 O concentrations are determined as follows. The optical depths (τν )
of the 3 µm stretching, 6 µm bending, and 13 µm libration modes of water are integrated over the frequency domain (ν in cm−1 ) and converted into the column density (N
in molecules cm−2 ), using the well known H2 O band strength values Aband [Hudgins et al.
1993], via:
R
τν dν
.
(4.2)
N=
Aband
The adopted H2 O band strength values, Aband , are 2 × 10−16 , 1.2 × 10−17 and 3.1 ×
10−17 cm molecule−1 for the stretching, bending and libration mode, respectively. The
H2 O column density is determined by taking the average of these three strong H2 O bands.
Similarly, at least four strong bands of the PAHs under investigation are integrated and
converted to column densities using the band strengths for these modes as calculated
above (Eq. 4.1). The average PAH column density is used to determine the concentration,
which is given by the ratio of the PAH and H2 O column densities.
Comparison of the PAH:H2 O ice spectrum before photolysis to that measured after
photolysis permits identification of photoproduct features, including cation absorption
83
4 IR spectroscopy of PAH containing ices
bands. The spectrum of the freshly deposited PAH:H2 O ice, multiplied by a factor, Y
(0 ≤ Y ≤ 1), is used as a reference that is subtracted from the spectrum of the irradiated
ice using the Resolutions Pro spectrometer software package provided by Biorad. The
factor Y is varied until the neutral bands are removed from the spectrum. At this point,
the resulting subtraction spectrum reveals only the photoproduct bands and the subtraction factor, Y, directly reflects the amount of neutral PAH consumed during photolysis.
Additionally, the contributions from the H2 O librational overtone and H2 O bending mode
are subtracted by a spline function, for which points are chosen at positions where no
PAH photoproduct absorptions are observed. The resulting baseline corrected spectra are
used for further analysis.
The sample window on which the ices are grown was thoroughly cleaned before commencing experiments on a different PAH. Thus, the background spectrum taken of the
cold sample window before starting a series of measurements (I0 ) was free of PAH absorptions. For one specific PAH, typically five individual PAH:H2 O photolysis experiments were performed for different concentrations and temperatures (Table 4.1), during
which a residue built up, comprising unprocessed PAH, and presumably also PAH:H2 O
photoproducts. After completion of the measurement series for this specific PAH, a spectrum (I) was taken of the room temperature “dirty” sample window. Subsequently, the
system was cleaned and prepared for the next run. The ratio ln(I/I0 ) reflects the optical
depth spectrum of the non-volatile residue and is used to derive complementary information on the species formed in the ice. The non-volatile residues that built up during the
Ant, Py, and Bghi P measurement series are discussed in §4.5.
4.3 PAH:H2 O spectroscopy
Earlier observations [Smith et al. 1989, Sellgren et al. 1995, Brooke et al. 1999, Chiar
et al. 2000, Bregman et al. 2000] indicate that the ratio of the number of PAH molecules
to the number of H2 O ice molecules is small along lines of sight towards protostars in
dense clouds. Based on these observations, we deduce that this number seems to be on
the order of a few percent. In a very careful study, Sellgren et al. [1995] reported the
optical depth of the 3.25 µm aromatic C–H stretch band towards Mon R2/IRS 3 as 0.045
with a FWHM of 75 cm−1 . These values are similar to the range of values reported by
Bregman & Temi [2001] towards other deeply embedded protostars. For the purposes of
this analysis, we use the results from Sellgren et al. [1995]. Similar conclusions can be
drawn using the observations from these other lines of sight.
Using the standard equation to determine the column density of absorbers along a
given line of sight:
τ × FWHM
,
(4.3)
N=
A
with the FWHM in cm−1 and an A value of 2.5 × 10−18 cm per aromatic C–H bond [e.g.
Joblin et al. 1994, Bauschlicher et al. 2008] yields 1.3 × 1018 aromatic C–H groups per
cm2 along the line of sight to Mon R2/IRS 3. Astronomical PAHs are thought to range in
size from roughly C25 to well over C100 . Reasonable formulae for such sized species are
84
W avelength /
3
m
4
6
0.7
8
6.5
7
7.5
8
8.5
9
12
9.5
16
20
10
Absorbance (-)
0.6
0.5
0.4
0.3
0.2
0.1
1600
0.0
1500
1400
1300
H O stretch
1100
1000
H O bend
2
3500
1200
H O libration
2
3000
2500
2000
2
1000
500
-1
Figure 4.1 A typical absorption spectrum for a pyrene:H2 O (1:90) ice mixture at 15 K. The inset shows a blow-up of that part of
the spectrum (region of interest) where Py absorptions are least affected by the water modes. The water modes are labeled in the
spectrum.
85
4.3 PAH:H2 O spectroscopy
Frequency / cm
1500
4 IR spectroscopy of PAH containing ices
C32 H15 and C130 H28 . To take this into account, we estimate that the ‘average’ number
of aromatic C–H bonds per astronomical PAH is 25. Dividing the column density of
aromatic C–H bonds towards Mon R2/IRS 3 (1.3 × 1018 ) by 25 yields the PAH column
density, 7.2 × 1016 PAHs cm−2 . Likewise, for Mon R2/IRS 3 we derive 3.3 × 1018 H2 O ice
molecules cm−2 from the O–H stretch band in Smith et al. [1989]. Thus, the PAH/H2 O
ice ratio along this line of sight is 7.2 × 1016 /3.3 × 1018 = 0.022, or 2%.
Figure 4.1 presents the 3750–500 cm−1 spectrum of a Py:H2 O (1:90) ice at 15 K,
which clearly shows that the H2 O ice bands dominate the spectrum at these concentrations. In agreement with earlier findings on other PAHs at concentrations of a few percent
in H2 O ice, the CH–stretch band near 3030 cm−1 (3.3 µm) is nearly imperceptible, while
the strong C–H out-of-plane bending bands between 900–500 cm−1 (11–20 µm) suffer
from severe blending with the H2 O ice libration mode [Sandford et al. 2004, Bernstein
et al. 2005a]. This makes PAH bands between 1650–1000 cm−1 (∼6.06–10 µm, the ROI)
the most promising to study PAH:H2 O photochemistry, and to identify astronomical PAHs
and PAH related photoproduct species in interstellar ices.
The spectroscopy of PAH:H2 O ice samples is discussed for Ant, Py, and Bghi P and
compared to available PAH matrix isolation data. In Fig. 4.2 typical baseline corrected
spectra of the three neutral PAHs in argon (trace A) are plotted together with the spectra of the neutral PAH:H2 O (trace B) ice (∼1:60). The spectra are normalized to the
strongest absorptions in the ROI. Clearly, the absorption bands are much broader in water
ice compared to the argon matrix data, causing some of the absorption bands to overlap.
Furthermore, the relative intensities of the bands differ for an argon or H2 O environment.
The band positions and integrated absorbances relative to that of the strongest PAH absorption are listed in Table 4.2 for Ant, Py and Bghi P in H2 O ice and in an argon matrix
[Hudgins & Sandford 1998a]. The corresponding Density Functional Theory (DFT) peak
positions taken from Langhoff [1996] are given as well. The absolute band strengths for
PAHs in H2 O ice, calculated under the assumptions described in Sect. 7.2, are also listed
in Table 4.2. Within the errors, the FWHM, relative intensity, and position of the band
maximum of absorptions of the three species studied here are found to be independent of
concentration and temperature for the values listed in Table 4.1, even for the two most
extreme Ant:H2 O ice concentrations (1:11 and 1:450). Thus, although we cannot fully
exclude the presence of PAH aggregates in the ice, there is no spectral evidence for such
species.
4.4 PAH ice photochemistry
All PAH:H2 O ice samples listed in Table 4.1 are VUV irradiated for 5, 10, 15, 30, 60,
120, and 180 minutes. Mid-IR spectra are taken after each VUV dose to look for changes
in the spectrum. Spline baseline corrected spectra after 5 minutes of VUV irradiation
of the Ant, Py, and Bghi P containing H2 O ice samples with a mixing ratio of ∼1:60 are
shown in trace (D) of Fig. 4.2. In this figure, the spectra are compared to those of the
corresponding PAH+ species produced and trapped in solid argon (trace C). Comparison
of the traces (D) and (B) in Fig. 4.2 shows that new absorptions arise upon photolysis of
86
4.4 PAH ice photochemistry
Table 4.2 The band positions and relative integrated absorbances for the PAHs antracene
(Ant), pyrene (Py) and benzo[ghi]perylene (Bghi P) in H2 O ice at 15 K, and in an argon
matrix at 10 K. Band positions computed using Density Functional Theory are also listed.
Absolute intensities (Aband ) for the PAHs in H2 O ice are in units 10−19 cm/molecule.
PAH
Ant
H2 O
1001.7
1103.8
1127.4
1149.7
1168.2
1186.8
1272.3
1316.0
1328.3
1347.7
1400.1
1450.9
1537.9
1563.1
Py
Bghi P
1624.4
1065.5
1096.4
1136.6
1176.3
1185.6
1244.0
1313.7
1435.1
1452.0
1468.4
1488.4
1594.1
1600.5
1038.3
1085.9
1149.2
Theorya
1000.7
1166.9
1156.2
1157.7
1169.3
1272.5
1318.1
1274.6
1311.2
1345.6,1346.4
1342.6
1450.5
1460.0
1542.0
1455.3
1456.1
1533.7
1610.5
1627.8
1620.0
1097.3
1092.3
1164.5
1183.9
1243.0
1312.1
1434.8
1160.8
1188.3
1253.1
1314.6
1427.0
1427.5
1471.0
1476.2
1132.0
1148.0
1186.5
1604.0
1036.3
1087.8
1093.4
1132.7
1149.2
1186.9
1203.5
1258.7
1206.7
1259.1
1305.6
1334.6
1342.2
1374.9
1302.9,1307.0
1392.8
1414.8
1447.8
1479.7
1512.8
1394.7
1416.0,1417.1
1449.0,1451.1
1582.8
1598.5
1614.7
1618.9
a Values
Position (cm−1 )
Ar
1000.9
1342.6
1350.0
1517.8
1527.4
1602.1
1586.1
1597.0
1037.6
1082.2
1091.9
1136.6
1152.7
1171.1
1206.6
1203.3
1261.5
1309.0
1292.3
1338.6
1336.0
1375.8
1398.9
1397.8
1426.2
1441.3
1513.3
1586.4
1586.6
Relative I.A.
H2 O
Ar
0.7
0.8
0.1
0.2
0.9
0.7
0.6
0.1
0.5
1.0
0.2
0.3
0.3
1.4
0.8
0.1
0.8
0.1
0.2
0.1
0.3
1.0
0.5
0.2
0.8
0.1
0.1
0.2
0.6
0.6
0.2
0.2
0.5
0.6
1.0
0.2
0.4
0.3
0.8
0.2
0.4
0.2
0.1
1.0
0.2
0.2
0.6
0.1
Aband
H2 O
6.8
0.8
1.9
9.0
6.2
1.0
5.7
10.4
2.1
3.5
3.1
14.6
8.7
0.8
7.9
2.8
5.1
2.1
6.1
21.7
10.3
5.3
17.3
2.8
1.8
3.3
12.7
12.0
4.4
3.6
0.4
0.8
0.3
0.4
0.1
0.0
0.0
0.3
0.1
0.0
0.5
0.1
0.3
0.0
10.8
1.7
0.7
0.3
1.0
0.1
0.3
1.0
0.2
15.8
6.3
22.5
1.6
0.5
0.5
0.8
0.2
0.2
0.4
0.3
0.1
0.1
0.4
0.3
0.6
0.0
0.1
0.2
8.5
17.8
6.2
11.7
11.4
17.8
4.7
3.5
9.1
6.5
1.3
2.4
taken from: Langhoff [1996]
87
4 IR spectroscopy of PAH containing ices
the PAH:H2 O mixtures. The photoproduct features are, however, heavily blended with
those of the remaining neutral PAH. To overcome this, spectra of freshly deposited PAHs
are subtracted from their VUV photolyzed counterparts with a subtraction factor, Y, as
described in §4.2. Subsequently, the contribution by the librational overtone and the H2 O
bending mode are removed by fitting and subtracting a spline function. The resulting
spectra are shown in trace (E) of Fig. 4.2 and show photoproducts only. The spectra
are again scaled to the strongest absorption in the ROI. It has to be noted here, that the
absorption bands of the photoproducts after 5 minutes of VUV irradiation are rather small
compared to those of the neutral bands, as can be seen by comparing trace (B) and (D)
from Fig. 4.2. While band intensities grow with longer irradiation time, the combination
of spectral overlap and different photoproduct band growths and losses makes it difficult
to track the precise photochemistry at longer photolysis times. This is discussed later in
§4.4.2.
It is clear from Table 4.2 that the relative PAH band strengths in H2 O ice are different
from those in an argon matrix, but the PAH cation band strengths measured in argon
have to be used for further analysis, since band strength values for PAH cation vibrational
modes in H2 O ice cannot be reliably determined for two reasons. First, the appearance of
bands other than cation absorptions after 5 minutes of photolysis points out that there is
no one-to-one conversion of the neutral PAH to the cation and, thus, the column density
of used up neutral PAH molecules cannot be used as a reliable column density for cations
in the ice on which to base a further band strength analysis. Secondly, spectral congestion
and ill-defined baselines make it hard to obtain accurate band areas of photoproduct bands
in irradiated mixtures. This choice induces additional errors in the absolute PAH+ column
density that can be as high as a factor of two.
4.4.1 PAH:H2 O photoproducts
VUV photoprocessed PAH containing H2 O ices spectra exhibit a set of broad absorption
features (see trace E of Fig. 4.2). As with the neutral PAHs, the absorption profiles of
PAH photoproducts in H2 O ice are much broader than those in an argon matrix (trace C).
These PAH:H2 O photoproduct absorption bands are decomposed by multiple Gaussian
fits and the peak positions of these bands are listed together with PAH+ absorptions and
band strengths measured in argon in Table 4.3. The solid lines along the top margin
indicate the PAH cation band positions in Ar [Hudgins & Allamandola 1995a,b]. The
photoproduct bands that fall within 10 cm−1 of these features are assigned to the PAH+
bands in H2 O. Some of the newly formed absorption bands, however, occur at positions
where no corresponding PAH+ band is found in argon. These absorption bands reflect
additional chemical reactions already at play in the early photolysis of PAH:H2 O ices.
It is well known that H2 O molecules photodissociate into radicals (H+OH) and that
these radicals are mobile within H2 O ice, even at low temperatures [Andersson & van
Dishoeck 2008, Öberg et al. 2009d]. At the concentrations under consideration here, it
is therefore likely that these and other photoproducts react with the PAHs, forming more
complex aromatic species containing functional groups that give rise to different peak
88
Wavelength
6.5
7
7.5
8
8.5
9
9.5
6.5
7
/
7.5
m
8
8.5
9
9.5
6.5
7
7.5
8
8.5
9
9.5
E
Absorbance (-)
D
C
B
A
1600
1500
1400
1300
1200
1100
1600
1500
1400
1300
Frequency / cm
1200
1100
1600
1500
1400
1300
1200
1100
-1
4.4 PAH ice photochemistry
89
Figure 4.2 From left to right the 1640 to 1050 cm−1 spectra of the PAHs anthracene, pyrene, and benzo[ghi]perylene considered
are shown here. A) Spectra of the neutral PAH in argon, B) spectra of the neutral PAH in water ice, C) spectra of the PAH cation
in argon, D) spectra of the PAH:H2 O (1:60) mixture after 5 minutes of in-situ VUV photolysis, and E) spectra showing only the
photoproduct features that appear after 5 minutes of in-situ VUV photolysis of the PAH:H2 O ice (spectrum ‘D’−YבB’). The tick
marks connected to the top axis indicate the positions of PAH+ features measured in argon. The PAH:matrix concentrations and
temperatures for the spectra shown in A and B are: Ant:Ar < 1 : 1000, 10 K; Py:Ar < 1 : 1000, 10 K; Bghi P:Ar < 1 : 1000, 10 K and
Ant:H2 O=1:60, 15 K; Py:H2 O=1:65, 15 K; Bghi P:H2 O=1:60, 15 K. Argon matrix spectra of neutral Ant and Py are reproduced from
Hudgins & Sandford [1998a], neutral Bghi P from Hudgins & Sandford [1998b], ionized Ant from Hudgins & Allamandola [1995b],
and ionized Py and Bghi P from Hudgins & Allamandola [1995a].
4 IR spectroscopy of PAH containing ices
positions in the mid-IR spectra (see also Chapter 6 and the discussion in Gudipati & Allamandola [2006b]). While infrared spectroscopic data on PAH:H2 O photoproducts are
largely lacking in the literature, analysis of the non-volatile photoproducts has shown
that the O, OH and H additions to the parent PAH are the dominant reaction pathways
[Ashbourn et al. 2007, Bernstein et al. 1999, 2002b]. Aromatic alcohols are among the
known photoproducts, and the C–O stretching vibration in alcohols and phenols produces
a strong band in the 1260–1000 cm−1 region [e.g., Silverstein & Bassler 1967]. Additionally, the alcohol OH wag falls in the 1420 to 1330 cm−1 region. Besides alcohols,
aromatic ketones are also amongst the photoproducts. The C=O stretch vibration in ketones typically occurs at around 1700 cm−1 . Keeping this in mind, tentative identifications
of unassigned bands, i.e. not due to PAH+ , are made below.
Anthracene:H2 O photoproducts
While most of the bands in the spectrum of the Ant:H2 O photoproducts can be attributed
to Ant+ , some prominent bands cannot. Two such bands appear at 1156 and 1243 cm−1 .
It is possible that both of these absorptions originate from the CO stretch of two different
Ant–OH isomers. The 1243 cm−1 band was previously attributed to an unknown Ant:H2 O
photoproduct [Bernstein et al. 2007]. Another prominent band that is not solely due to
Ant+ occurs at 1452 cm−1 . While a small cation band is present near this position in the
argon matrix data, it is much less intense with respect to the other cation bands whereas
the photoproduct band in trace (E) (Fig. 4.2) is one of the strongest in the spectrum.
This band is likely a blend of the cation band with a much stronger product band. The
absorption frequency suggests that it originates from an aromatic CC stretching and C–H
in-plane bending mode. Two strong photoproduct bands also appear in the blue end of
the ROI, one at 1518 cm−1 , the other at 1604 cm−1 . The moderately strong Ant+ band at
1586.4 cm−1 is expected to contribute to the blue end of the 1604 cm−1 feature, but again
seems to be too weak to explain the full feature. Both the 1518 cm−1 and most of the
1604 cm−1 bands are likely due to the aromatic CC stretch vibration of a newly formed
species.
The detection of new bands in the 1260–1000 cm−1 region upon VUV photolysis of an
Ant:H2 O mixture does agree with earlier findings by Ashbourn et al. [2007], who detected
1-anthrol and 2-anthrol using HPLC in a VUV irradiated Ant:H2 O ice (> 1:100) after
warm-up to room temperature. They also reported the formation of 1,4-anthraquinone,
9,10-anthraquinone and 9-anthrone. Anthraquinones contain two C=O bonds, which typically absorb at a frequency of 1676 cm−1 [Chumbalov et al. 1967]. This band position is
outside of our ROI and detection is hampered by the strong H2 O bending mode. The other
species detected by Ashbourn et al. [2007], 9-anthrone, belongs to the group of ketones
and is expected to exhibit a C=O absorption at ∼1700 cm−1 . The formation of this species
can also not be confirmed for the same reason.
90
4.4 PAH ice photochemistry
Table 4.3 Band positions of photoproducts appearing upon VUV photolysis of the PAHs
antracene (Ant), pyrene (Py) and benzo[ghi]perylene (Bghi P) in water ice at 15 K compared to cation absorption band positions and band strengths measured in an argon matrix.
Band strengths are in units of 10−18 cm/molecule. Cation bands are marked with a ’+’.
PAH
Ant
Position (cm−1 )
H2 O
1155.7
1190.5
1242.7
1292.8
1326.2
1340.5
1358.3
1412.1
Py
Bghi P
1419.5
1424.7
1452.2
1459.6
1517.9
1541.0
1567.2
1589.3
1603.6
1621.1
1137.3
1182.0
1205.2
1216.7
1232.2
1247.6
1295.4
1319.0
1337.7
1359.0
1372.8
1393.2
1422.2
1446.0
1484.0
1537.5
1553.4
1567.4
1586.4
1614.0
1080.2
1126.4
1146.7
1191.0
1223.6
1253.1
1303.0
1322.9
1339.8
1346.2
1352.9
1366.8
1380.3
1388.4
1406.1
1433.3
1479.7
1501.0
1510.4
1529.7
1550.5
1568.0
1578.0
1590.2
1604.7
1617.0
1624.3
a Values
Positiona (cm−1 )
Argon
Band strengtha
Ass.
1183.3
1188.6
0.37
18
+
+
1290.4
1314.6
1341
1352.6
1364.4
1406.1
1409.5
1418.4
1430.2
1.5
1.5
26
8.0
1.0
0.39
2.7
22
0.38
+
+
+
+
+
+
+
+
+
1456.5
1.9
+
1539.9
3.9
1586.4
3.6
+
+
+
1188.7
0.56
+
1216.0
1.7
+
1245.1, 1253.7,1255.7
5.6
+
1356.1,1358.4, 1361.8
16
+
1421.1
1440.3
2.3
1.5
+
+
1550.9,1553.4,1556.0
13
+
1140.2
5.5
+
1216.7
1223.4
0.46
4.1
+
+
1311.9
1324.4
1331.9
1.9
10
2.3
+
+
+
1350.2
1369.0
2.1
8.6
+
+
1388.3
1401.3
1408.8
1429.4
0.44
12
0.64
1.3
+
+
+
+
1538.6
1550.1
0.25
3.1
+
+
1578.2
14
+
taken from: Hudgins & Allamandola [1995a,b]
91
4 IR spectroscopy of PAH containing ices
Pyrene:H2 O photoproducts
In the Py:H2 O case, the spectra in Fig. 4.2, trace (E) and photoproduct peak positions
listed in Table 4.3 show that less than half of the new bands can be confidently assigned
to Py+ . As discussed above for Ant, the new features between 1260 and 1000 cm−1 may
be due to the CO stretch in various Py–OH isomers. Likewise, the prominent features
at 1373 and 1393 cm−1 can be tentatively attributed to modes involving both aromatic
CC stretch and C–H in-plane bends of new products. The large number of photoproduct
bands between 1500 and 1650 cm−1 is striking, especially since only that at 1553 cm−1
can be confidently attributed to Py+ based on matrix isolation spectra.
An attempt has been made to assign the unknown absorptions to more complex Py
related species. Peak positions in the Py:H2 O photolysis experiments are compared to
peak positions of 43 pyrene related species from the extensive theoretical database of
PAH derivatives [Bauschlicher et al. 2010,www.astrochem.org/pahdb]. Some groups of
molecules indeed do exhibit strong transitions around the peak positions where the absorption maxima are found for the undefined Py:H2 O photoproducts. These molecules
include H, OH, and O added pyrene-based species, such as C16 H11 O, C16 H10 O, C16 H12 ,
and their cations, in a variety of possible configurations. Although some of these theoretical peak positions overlap with the photoproduct bands, accurate experimental spectral
data for these molecules are needed for unambiguous identifications of the reaction products of Py in VUV photolyzed H2 O ice.
Benzo[ghi]perylene:H2 O photoproducts
As molecular size increases, the number of mid-IR transitions grows and, because of
spectral congestion, subtraction of neutral precursor bands becomes increasingly difficult.
This makes it hard to obtain clear-cut spectra of the products in the Bghi P photoprocessed
ice and hence makes identification of individual bands difficult, if not impossible. As
with Ant and Py, the spectra in Fig. 4.2 and photoproducts listed in Table 4.3 show that
several absorption bands appear which clearly do not have a matrix cation counterband.
Unassigned absorptions are found between 1470 and 1540 cm−1 and probably involve CC
stretching and C–H in plane bending modes. The strong unidentified bands between 1600
and 1640 cm−1 are likely caused by the C=O stretching mode of Bghi P ketones formed in
the ice.
Keeping the types of photoproducts in the Ant and Py ices in mind, it is most likely
that Bghi P derivatives containing H, O, and OH groups are formed upon photolysis. As
for the Ant:H2 O and Py:H2 O experiments, we cannot unambiguously assign the Bghi P
photoproduct absorption bands. The non-volatile residue of a VUV irradiated Bghi P:H2 O
(<1:800) ice shows the addition of O, OH and H to the neutral parent [Bernstein et al.
1999]. Thus, as with the other PAH:H2 O systems studied to date, it is likely that many
of these new bands in the mid-IR are due to various forms of Bghi P–OH, Bghi P–O, and
Bghi P–Hn and possibly their ionized counterparts.
92
4.4 PAH ice photochemistry
4.4.2 Concentration effects and time dependent chemistry
PAH:H2 O photolysis experiments have been performed for a set of concentrations ranging
from ∼ 1:11 to 1:200. Here only the Ant:H2 O experiments are described, but all three
investigated PAHs exhibit a similar behavior. Figure 4.3 shows the decay in the amount
of the neutral parent Ant in the ice as a function of photolysis time relative to the amount
of the freshly deposited Ant before irradiation. Clearly, Ant loss is far more efficient for
lower than for higher PAH concentration. Extrapolating these results, they are in good
agreement with recent results that are described in Chapter 7 for PAH:H2 O ices at very
low concentration (1:∼5,000 to 10,000), where it is found that all the neutral PAH was
consumed at the end of a 4 hour photolysis experiment.
VUV fluence / 10
Remaining fraction of neutral anthracene (-)
0
20
40
17
photons cm
60
-2
80
100
1.0
0.9
0.8
0.7
0.6
0.5
0.4
0.3
Ant:H O
0.2
2
1:11
0.1
1:60
1:172
0.0
0
20
40
60
80
100
120
140
160
180
Photolysis time / minutes
Figure 4.3 Neutral anthracene decay as a function of photolysis time (VUV fluence) and
concentration for 1:11, 1:60, and 1:172 Ant:H2 O mixtures at 15 K. Conservative errors
are ±5, ±7, and ±10%, respectively, of the initial amount of deposited neutral species.
Second order exponential fits to the data (solid lines) are shown as well.
The 15 K neutral Ant decay data in Fig. 4.3 are co-plotted with a fit of the form
Ψ=C1 exp(−t/τ1 ) + C2 exp(−t/τ2 ). Fitting the experimental data required a sum of two
exponentials. This indicates that more than one process is responsible for the neutral Ant
loss. This is consistent with the optical work presented in Chapters 6 and 7 in which two
photochemical reaction networks were described, one for PAH cation formation, the other
involving H, O, and OH PAH addition reactions.
93
4 IR spectroscopy of PAH containing ices
The time dependent PAH cation and other photoproduct signals are also studied as
a function of photolysis time (VUV dose). For the lowest concentration Py:H2 O ice in
our sample (1:200), the time dependent behavior of the photoproduct bands is compared
to the optical results presented in Paper II. To this end, the behavior of the cation is
traced as a function of VUV time by integrating two of the most isolated prominent Py+
absorption bands in the spectra, located at 1359 and 1554 cm−1 . Figure 4.4 shows the time
evolution of the column density of the Py+ species derived from these bands (multiplied
by a factor of 10 to facilitate the display and normalized to the initial amount of neutral Py)
together with the time evolution of the relative amount of neutral Py in the ice based on the
subtraction factor, Y. The photolytic behavior of the integrated absorption of the strongest
undefined photoproduct band at 1567 cm−1 (multiplied by a factor of 100 to facilitate the
display) is also shown. Because no information is available on the band strength of this
species, we cannot convert this integrated absorbance into a column density relative to the
amount of deposited neutral PAH. The Py+ bands and 1567 cm−1 photoproduct band show
a different time dependence and clearly do not correlate. While the integrated absorbance
of Py+ reaches a maximum after some 10 minutes and then declines, the photoproduct
signal grows and levels slightly off towards the end of the experiment.
VUV fluence / 10
0
20
40
17
photons cm
60
-2
80
100
1.0
Py
0.9
+
Normalized column density
Py
0.8
x 10
Product
0.7
0.6
0.5
0.4
0.3
0.2
0.1
0.0
0
20
40
60
80
100
120
140
160
180
Photolysis time / minutes
Figure 4.4 Evolution of Py, Py+ and a single Py photoproduct for a 1:200 Py:H2 O ice at
15 K as a function of photolysis time (VUV fluence). Photoproduct species are tracked
using the 1359 and 1553 cm−1 Py+ and 1567 cm−1 Py photoproduct bands.
The error in the Py+ column density is reduced by taking the average column density
based on the two strongest cation absorptions located at 1359 and 1554 cm−1 . We estimate
the error in the relative Py+ abundance to be ±5% of the deposited amount of neutral Py
94
4.4 PAH ice photochemistry
and, thus, the shape of the time dependence curve is well defined. The error in the absolute
number, however, is much larger because the argon band strength values from Table 4.3
have to be used for the analysis and because of ill-defined baselines caused by spectral
congestion. Therefore, the error in the absolute amount of Py+ produced can be quite
high. For this reason, reliable quantitative time dependent data can only be obtained in
the optical spectral regime, where cation absorptions are isolated and spectra are taken
on a much shorter time scale. This is described in Chapter 7. The error in the absolute
number of neutral Py molecules consumed, on the other hand, is low (±7%). Since no
information is available on the band strength of the 1567 cm−1 unidentified photoproduct
absorption band, no quantitative information can be obtained on the formation of the
species responsible for this band.
Using the PAH cation in Ar band strengths values, it is also possible to derive Nfrac ,
the first order estimate of the fraction of the used-up neutral PAH that is converted into
cation species versus other photoproducts. The amount of deposited neutral, NPAH , is
known from concentration measurements (Eqns. 4.1 and 4.2) and the amount of consumed
neutral is given by (1 − Y)NPAH . The cation column density, NPAH + , is calculated with
Eq. 4.2 and the band strength values reported in Table 4.3. The difference between the
+
two, (1 − Y)NPAH − NPAH
, makes up the photoproduct column density. The fraction of
consumed neutral PAH molecules converted into PAH+ is given by:
Nfrac =
NPAH+
.
[(1 − Y)NPAH ]
(4.4)
Values Nfrac have been derived for different conditions for the three PAHs studied
here. After 5 minutes of VUV photolysis of the Py:H2 O (1:200) mixture, roughly 15%
of the used up neutral Py is converted into Py+ and 85% in other charged and/or neutral
photoproducts. The Py+ absorption peaks at about 10 minutes of VUV photolysis, a
value that is also found in the optical study for much lower PAH concentrations (1:5,0001:10,000). A mid-IR experiment on an ice sample with a higher Py concentration (1:60)
results in an even lower Nfrac value of about 0.05, a trend that suggests that ionization is
more important in low concentration ices. Similar trends are found for low concentration
mixtures of Bghi P — a (1:160) mixture gives Nfrac ≈ 0.20 — and Ant — a (1:172) mixture
gives Nfrac ≈ 0.10 — after 5 minutes of VUV photolysis.
The quantitative analysis in Chapter 6 points out that there is a Py+ +e− recombination
channel. It is argued that the recombination reaction is most likely a local process, i.e.,
the electron released after ionization remains in the vicinity of its parent Py molecule.
Putting this in perspective with the data presented here, this means that the recombination
channel can well be more efficient in ices of higher concentration; the chance of recombination with an electron released from a neighboring Py molecule is larger. In the higher
concentration experiments presented here, not all the neutral PAH is used up. Presumably
some of the deposited PAHs are shielded from VUV irradiation and therefore not ionized.
This may also explain the difference between the ionized fraction in low versus high concentration ices. Furthermore, as discussed before, ionization seems to be less efficient at
high concentration ices, because chemical reactions between PAHs and radical species
likely dominate the loss of the neutral species.
95
4 IR spectroscopy of PAH containing ices
In summary, all PAH molecules trapped in H2 O ice exhibit a similar photoprocessing
behavior. Ionization is more important in low concentration ices. However, after 5 minutes of VUV photolysis, most of the destroyed neutral species are converted into different
PAH based photoproducts other than the cation. The concentration of the ice sample has
a large influence on the efficiency of the chemical reactions. PAHs in lower concentration ices are used up faster and more efficiently, whereas in higher concentration ices,
the PAH consumption is far less efficient. Although no mid-IR VUV photolysis measurements have been performed on coronene:H2 O ice mixtures, the near-UV/Vis study
presented in Chapter 7 shows that also coronene exhibits a similar behavior.
4.4.3 Ionization efficiency in CO ice
A control experiment has been performed on a Py:CO ice sample at 15 K to investigate the
ionization efficiency of PAHs upon photolysis in a CO matrix. After a short photolysis
time (150 s), the ice exhibits weak absorptions at 1861 and 1090 cm−1 . These midIR absorption bands were previously assigned to the HCO· radical by Milligan & Jacox
[1969]. This observation is consistent with experiments reported in Chapter 6, where
the electronic HCO· signature is found in the 500 to 660 nm spectral range. There, it
is also found that, in a CO matrix, the Py ionization efficiency is close to zero, unless a
certain level of H2 O contamination in the CO matrix is reached. In the mid-IR experiment
performed here, VUV photolysis of Py in a CO matrix indeed does not show any sign of
pyrene ionization, but also does not show any other PAH:H2 O photoproduct bands. We do
confirm the low ionization efficiency and the appearance of small HCO· absorptions in a
nearly pure CO ice, but no absorptions caused by PyH· species, as in the optical study, are
observed. The fact that PAH absorption band strengths typical of electronic transitions
are 100 times stronger than those for vibrational transitions, and that the level of H2 O
contamination, i.e. the source of H-atoms for the reaction H· + CO· →HCO· , is lower in
the experiments described here, probably explains why the PyH· mid IR absorption bands
were not detected. The important conclusion that follows from this control experiment
is that H2 O catalyzes the ionization process. This is also consistent with the observation
that ionization seems to be more efficient in low concentration PAH:H2 O ices.
4.4.4 Temperature effects
VUV photolysis experiments on PAH:H2 O samples (∼1:60) were also conducted at a
higher temperature (125 K). Consistent with the behavior reported in the optical studies
Chapter 6, we find that the neutral PAH loss is still efficient, while the ionization channel
is strongly suppressed. The less efficient formation of the cation can point to a lower rate
of ionization, but it is also possible that the recombination channel becomes dominant at
higher temperatures. The result is that, at higher temperatures, the parent PAHs are more
efficiently converted into species other than the PAH cation.
In the high temperature Py:H2 O experiment, for example, some pronounced vibrational bands appear immediately upon VUV photolysis. These bands, located at 1137.8,
96
4.5 The non-volatile residue
1216.7, 1386.1 cm−1 and a broad feature consisting of bands at 1553.5, 1566.1, and
1583.1 cm−1 , become strongest after 5 minutes of photolysis and then subside. Only
the bands at 1216.7 and 1553.5 cm−1 are at positions of Py+ absorptions. The remaining bands are also apparent in the low temperature spectra, but shifted by up to 3 cm−1
and, in the absence of many of the other bands, seem to be more pronounced in the high
temperature photolysis dataset.
For the PAH:H2 O ices irradiated at high temperatures there is a very broad overlapping substructure superimposed on the baseline of their mid-IR spectra. This is presumably caused by blended photoproducts and possibly PAH aggregates. Together with the
knowledge that more complex species have been detected in other experiments [Bernstein
et al. 1999, 2002b, Ashbourn et al. 2007], we conclude that predominantly a mixture of
PAH–Xn species, with X being H, O, or OH, may have been formed and that only one or
a few chemical reaction channels dominate, resulting in the observed bands.
Ices at a temperature of 125 K are known to be of different structure than ices at low
temperature [Jenniskens & Blake 1994]. This structural difference — amorphous at 15 K
vs. cubic crystalline at 125 K — may explain the different photochemical behavior. However, in Chapter 6 we investigated the influence of the ice structure on the photoionization
and found that ices annealed at 125 K and subsequently cooled down to 25 K exhibit similar ionization behavior as those grown and photolyzed at 25 K. Hence, they ascribed the
different behavior at high temperatures to the larger mobility of radical species in the ice.
The experiments presented here, point out a similar behavior.
4.5 The non-volatile residue
Figure 4.5 shows the 4000 to 500 cm−1 room temperature spectra of the non-volatile
residues produced by the photolysis of the three PAH:H2 O ices considered here. The
spectra are measured as described in §7.2 and provide additional information on the photoproducts which have accumulated over a series of experiments for a particular PAH.
These residues, complex mixtures of the non-volatile photoproducts, may also contain
some of the parent PAH as well as some trapped H2 O or H2 O that accreted during the
cooling of the sample window. Because the H2 O absorption bands in these spectra are
much smaller than in the experiments on PAH:H2 O ice mixtures, the spectra can be investigated over the full range from 4000 to 500 cm−1 . In the ROI, the spectra show
continuous, undulating absorptions from about 1750 to 1130 cm−1 with several distinct
features superposed. Additional spectral features are found between 3750 and 2750 cm−1 .
The chemical subgroups indicated by these features are consistent with the addition of O,
H, and OH to the parent PAH.
The aromatic-rich nature of these residues makes their spectra qualitatively different
from the 13 residue spectra previously analyzed to compare with the spectrum of the
diffuse interstellar medium [Pendleton & Allamandola 2002]. While it is impossible to
identify the molecules comprising the residues using infrared spectroscopy, it is possible
to characterize the species and subgroups present by chemical type (i.e aliphatic versus
aromatic hydrocarbons, carbonyl vs. alcohol carbon-oxygen links, etc.). The following
97
4 IR spectroscopy of PAH containing ices
analysis utilizes the characteristic group frequencies as well as relative and intrinsic band
strengths as described in [Bellamy 1960, Silverstein & Bassler 1967, Wexler 1967]. The
prominent features in these spectra are discussed systematically from higher to lower
frequency.
Wavelength /
3
m
5
10
15 20
Ant
Optical Depth (-)
Py
B
ghi
4000
P
3000
2000
Frequency / cm
1000
-1
Figure 4.5 Overview of the room temperature spectra of the residues accumulated during
the photolysis of Ant, Py, and Bghi P containing H2 O ices.
• The 3300 cm−1 , O–H stretch: A large part of this band can be ascribed to H2 O accretion to the cold sample, since the residue spectra are based on background spectra taken of a cold sample window. However, the weak sideband near 3550 cm−1
points to the O–H stretching mode in phenols and, thus, the presence of aromatic
alcohols in the residues.
• The 3060 cm−1 aromatic C–H stretch: The next prominent band, peaking near
3060 cm−1 , is due to the aromatic C–H stretch belonging to either the parent PAH
or PAH photoproducts.
• The 2990–2850 cm−1 aliphatic C–H stretch: The addition of H to the aromatic parent forms a new type of functional group, which is evident from the broad aliphatic
C–H stretch feature between about 2850 and 2990 cm−1 . The two subpeaks at
about 2925 and 2860 cm−1 are characteristic of methylene (–CH2 –) groups, showing that the major H-addition channel proceeds by addition, not ring opening. If
ring opening had been an important channel, bands due to the C–H stretches of
methyl (–CH3 ) groups near 2960 and 2870 cm−1 would have been expected. The
98
4.5 The non-volatile residue
assignment to methylene groups is further supported by the prominent band near
1450–1460 cm−1 in the lower two spectra shown in Fig. 4.5. A band in this region
indicates the –CH2 – scissoring mode in cyclic saturated compounds such as cyclohexane and cyclopentane. The spectrum of Bghi P is consistent with this analysis, but
the bands are weaker. Given that the intrinsic strength of the aromatic C–H stretch
is 10 times weaker than the intrinsic strength of the aliphatic C–H stretch per C–H
[Wexler 1967], the increasing ratio of the aromatic band intensity to the aliphatic
band intensity with PAH size indicates that H addition becomes less effective as the
PAH size increases.
• The 1700 cm−1 carbonyl C=O stretch: As with the O–H stretching band near 3 µm,
some of this band may be attributed to the bending mode of H2 O molecules frozen
onto the cold sample window. However, it is likely that part of this band is caused by
the C=O stretch in carbonyl groups. In this case, quinone-like structures would be
expected with the double bond to O participating in the conjugation of the molecule
as a whole.
• The band at 1625 cm−1 : A distinct feature is evident at this frequency in all three
spectra. The value is low for isolated carbonyls, but is characteristic for the carbonyl
C=O stretch in molecules with intramolecular hydrogen bonding with an OH group.
• The 1600 cm−1 aromatic CC stretch: A well defined, but generally weak band
centered near 1600 cm−1 is characteristic of the CC skeletal stretch in aromatic
molecules. The prominence of this band in all three residue spectra shows, not
surprisingly, that they are primarily aromatic in nature.
• The bands between 1520 and 1500 cm−1 : Several features in this region become
more prominent with parent PAH size. We are unable to attribute these bands to
specific molecular vibrations.
• The 1450 and 1440 cm−1 –CH2 – scissoring vibration: This band has been discussed
in the paragraph describing the aliphatic C–H vibrations above.
• The 1260 cm−1 band in the pyrene photoproduct: While this prominent band falls in
the 1300–1000 cm−1 region which is characteristic of aromatic C–H in-plane bending vibrations, other likely compounds also absorb in this region. For example, the
C–O stretch in aromatic alcohols has a strong band between 1260 and 1180 cm−1 .
• The 900 to 500 cm−1 aromatic C–H out-of-plane bends: This region encompasses
the very strong C–H out of plane vibrations which are diagnostic of substitution
patterns on the edge rings. In this case, these bands serve primarily as an indication
of the presence of aromatic compounds.
All the pure PAHs under investigation here exhibit strong fluorescence upon excitation with a hand-held UV lamp. It is noteworthy, that after performing the experiments,
the residues do not exhibit any fluorescence. This is consistent with the residues being
processed PAHs with a different chemical identity.
99
4 IR spectroscopy of PAH containing ices
4.6 Astrophysical implications
The 5–8 µm spectra of embedded low- and high-mass objects are generally dominated
by two strong bands originating from interstellar ice, one centered near 6.0 µm and primarily attributed to the H2 O bending mode, and the other near 6.8 µm. Close inspection
of the spectra of many different molecular clouds shows that, while features near 6.0 and
6.8 µm are nearly always present, band positions, profiles, and relative strengths of other
features vary somewhat from one object to another, revealing changes in ice composition and structure [e.g. Keane et al. 2001b, Keane 2001, Gibb & Whittet 2002, Schutte
& Khanna 2003, Boogert & Ehrenfreund 2004, Boogert et al. 2008]. Besides the two
dominant absorptions, signals are also found from other species that absorb in the 5–8 µm
region. An overview of absorptions by molecular species in this wavelength region is
given in Table 4.4. The knowledge of interstellar ice composition and chemistry is based
on over twenty five years of dedicated studies on laboratory ice analogs to interpret these
observations [see references mentioned above and Ehrenfreund & Charnley 2000, van
Dishoeck 2004,and references therein]. However, to date, such studies considering PAHs
in ices and their contribution to the 5–8 µm absorption complex are lacking. Here, we discuss the possible contribution of PAHs in ices to the spectra of both high- and low-mass
objects, using the laboratory data presented in the previous sections.
Table 4.4 Molecules, their vibrational modes and solid state absorption band positions
thought to contribute to the 5–8 µm ice absorption complex towards low- and high-mass
embedded objects
Mode
λ (µm)
ν (cm−1 )
Ref.
H–O–H bend
5.4–9
1852–1111
1
Libr. overtone
1
H2 CO
C=O stretch
5.83
1715
2
CH3 HCO C=O stretch
5.83
1715
3
HCOOH
C=O stretch
5.85
1709
3,4
NH3
H–N–H bend
6.16
1623
5
PAH?
CC stretch
6.20
1613
6
HCOO−
C=O stretch
6.33
1580
3
H2 CO
H–C–H bend
6.68
1497
2
NH+4
Deformation
6.85
1460
7
HCOOH
C–H stretch
7.25
1379
3,4
HCOO−
C–H deform.
7.40
1351
3
SO2
Asymm. stretch
7.58
1319
8
CH4
Deformation
7.68
1302
9
References: [1] D’Hendecourt & Allamandola [1986]; [2] Schutte et al. [1993]; [3] Schutte et al.
[1999]; [4] Bisschop et al. [2007a]; [5] Kerkhof et al. [1999]; [6] Keane et al. [2001b] Kerkhof et al.
[1999] ; [7] Schutte & Khanna [2003]; [8] Boogert et al. [1997]; [9] Boogert et al. [1996].
Molecule
H2 O
100
4.6 Astrophysical implications
4.6.1 High-mass protostars
Figure 4.6 shows the ∼5.3–8.7 µm (1900–1150 cm−1 ) ISO SWS spectrum towards the
high-mass protostar W33A (based on data from Keane et al. [2001b]) and the residual
that remains after subtracting the H2 O bending and librational overtone bands in a manner
similar to that described in §7.2. Based on a sample of five ISO spectra towards high-mass
protostars, Keane et al. [2001b] attributed the bulk of the 6.0 µm ice band to amorphous
H2 O ice, with additional absorptions on the short wavelength wing near 5.8 µm and on
the long wavelength wing near 6.2 µm. They assigned the short wavelength absorption
at 5.83 µm to the carbonyl CO stretch in formaldehyde and related species based on
laboratory studies and tentatively suggested that the long wavelength absorption at 6.2 µm
is due to the CC stretch of aromatic structures in PAHs or amorphous carbon particles.
Keane et al. [2001b] also showed that the 6.85 µm band is comprised of two components,
one centered near 6.75 µm, and the other at 6.95 µm, but they did not assign possible
band carriers. Noteworthy in Fig. 4.6 are also the weak undulating features between ∼7.2
and 7.9 µm. Assignments of these absorption features are discussed in Schutte & Khanna
[2003].
Wavelength /
5.5
6
m
6.5
7
7.5
8
8.5
-0.2
0.0
0.2
Optical Depth
0.4
0.6
0.8
1.0
1.2
1.4
W33A
1.6
W33A - H O
2
1.8
1800
1700
1600
1500
Frequency / cm
1400
1300
1200
-1
Figure 4.6 The 5.3 to 8.7 µm ISO SWS spectrum of the high-mass protostar W33A plotted
in optical depth together with the spectrum obtained after subtraction of the H2 O bending
and librational overtone modes. The figure is adapted, with permission, from Keane et al.
[2001b].
101
4 IR spectroscopy of PAH containing ices
4.6.2 Low-mass protostars
Figure 4.7 shows the ∼5.3–8.1 µm (1900–1234 cm−1 ) spectrum of the low-mass YSO
RNO 91, reproduced from Boogert et al. [2008]. In this more recent, extensive spectroscopic study of the 5–8 µm region of 41 low luminosity, low-mass protostars measured
with the Spitzer telescope, Boogert et al. [2008] showed that, once the contribution from
H2 O ice is removed, the 5–8 µm residual spectrum can be split into five components
(C1–C5). The large sample of objects permitted Boogert et al. to extract rough band profiles, although not uniquely assignable to specific band carriers, and to determine band to
band variations for all five components. The positions of the C1 through C4 components
are similar to those described in Keane et al. [2001b] and part of these bands can be attributed to the species mentioned in Table 4.4. C5, a very broad, underlying component
that stretches from about 1725 to 1250 cm−1 (5.8–8 µm) cannot be explained by any of
these species and is most likely a blend of more than one species. As with the spectrum
of W33A, RNO 91 also shows several weaker features between 7.2 and 7.9 µm.
Wavelength /
5.5
6
m
6.5
7
7.5
8
-0.05
0.00
0.05
Optical Depth
0.10
0.15
0.20
0.25
0.30
0.35
RNO 91
0.40
RNO 91 - H O
2
Component C5
0.45
1900
1800
1700
1600
1500
Frequency /cm
1400
1300
-1
Figure 4.7 The 5.3 to 8.1 µm Spitzer spectrum of the low-mass YSO RNO 91 plotted in
optical depth together with the spectrum obtained after subtraction of the H2 O bending,
and librational overtone modes. The C5 complex absorption as described by Boogert et al.
[2008] is also indicated (solid line). The figure is adapted, with permission, from Boogert
et al. [2008].
102
4.6 Astrophysical implications
4.6.3 PAH contributions to the 5–8 µm absorption
Here, we discuss the possible contribution that interstellar ices containing PAHs and their
photoproducts could make to the spectra of molecular clouds in the 6.1–8.3 µm (1650–
1200 cm−1 ) region. The comparison is restricted to the three PAHs studied here. These
three PAHs, Ant in particular, are on the low end of the PAH size distribution expected in
space. As a consequence, the relative intensities of the modes are different than those expected for larger PAHs in the interstellar case. Infrared bands which result from the C–H
functional group are expected to be less intense in space with respect to modes resulting from aromatic CC stretches, since astronomical PAHs are large and contain relatively
more CC than CH bonds. Figure 4.8 shows the photoproduct bands that appear after
180 minutes of VUV photolysis of the PAHs considered here in H2 O ice. All VUV processed PAH:H2 O mixtures under investigation exhibit bands clustered around 1600 cm−1 .
The band at 6.2 µm, observed towards both low- and high-mass sources and also shown
in Fig. 4.8, could well be caused by a superposition of many different large PAH and
Wavelength /
6.5
m
7
7.5
8
Ant prod.
Optical Depth (-)
Py prod.
BghiP prod.
RNO 91-H O
2
W33A-H O
2
1650
1600
1550
1500
1450
1400
Frequency / cm
1350
1300
1250
1200
-1
Figure 4.8 The 6.1–8.3 µm (1650 cm−1 –1200 cm−1 ) region of the spectra of W33A and
RNO 91 after subtraction of the H2 O bending and librational overtone contributions plotted in absorbance compared to the spectra of the products after 180 minutes of PAH:H2 O
ice photolysis. The dotted line guides the eye to the ∼6.2 µm ice absorption and main
PAH:H2 O ice photoproduct bands.
103
4 IR spectroscopy of PAH containing ices
PAH photoproduct absorption bands, or organic refractory material as discussed in §4.5.
Additional minor PAH absorptions throughout the 6–10 µm spectrum can be (partly) responsible for the underlying C5 complex absorption which is also indicated in Fig. 4.5 as
well as the weak undulating structure between 7.2 and 7.9 µm.
The observational data also permit one to put crude limits on the column densities
of PAHs condensed in interstellar ice. Assuming that the 6.2 µm feature in the residual
spectrum shown in Fig. 4.6 is due primarily to the aromatic CC stretching vibration in
PAHs and closely related aromatic material frozen in H2 O-rich interstellar ice, one can
estimate their column density (NPAH ) with Eq. 4.2. The integrated optical depths of the
component are ∼16 cm−1 for W33A and ∼1.3 cm−1 for RNO 91. The experimental data
indicate that a certain number of neutral PAH species is converted into PAH photoproduct
species, as indicated by (1 − Y)Nneutral , where Nneutral is the column density of deposited
neutral molecules. This allows us to calculate an average band strength for the species responsible for the absorption band around 6.2 µm with Eq. 4.2. The resulting average band
strength (Aprod ) for the formed PAH products is 1.2 × 10−17 cm molecule−1 . With these assumptions, the column density of neutral PAHs and PAH photoproducts frozen in H2 O is
1.3×1018 molecules cm−2 for the high-mass object (W33A) and 7.6×1016 molecules cm−2
for the low-mass object (RNO 91). Given that the column densitiy of H2 O ice along the
line of sight to these objects is 4 × 1019 molecules/cm2 and 4.3 × 1018 molecules/cm2
[Keane et al. 2001b, Boogert et al. 2008], respectively, this implies PAH and PAH photoproduct abundances between 2–3% with respect to water. This concentration range is
reasonable given the PAH abundances derived from mid-IR emission features [Puget &
Leger 1989, Allamandola et al. 1989].
4.7 Conclusions
The manuscript describes mid-IR absorption spectroscopy of PAH:H2 O ices, their VUV
induced photochemistry and photoproducts. The main conclusions of this manuscript are
summarized below.
1. Band positions, FWHMs, and relative intensities of anthracene, pyrene and benzo[ghi]perylene in H2 O ice in the 1650–1000 cm−1 spectral window are compared
to their matrix isolation data. Additionally, band strengths are derived for these
species isolated in an astrophysically relevant H2 O ice at 15 K.
2. VUV photolysis of PAH containing H2 O ice causes the embedded PAH to ionize
and react with H2 O photoproducts. The band positions of PAH cations trapped in
H2 O ice are measured and compared to PAH cation values from previous argon
matrix isolation studies. Additionally, peak positions of PAH:H2 O photoproducts
other than the cation are determined and tentatively assigned to PAH–Xn species,
with X being O, H or OH.
3. The PAH:H2 O ice photochemistry is tracked as a function of VUV fluence by monitoring the integrated absorbances of the parent PAH and PAH photoproduct bands
104
4.7 Conclusions
periodically during photolysis for several hours. In all cases, the PAH cation bands
peak after some 5 to 10 minutes and then slowly decline while other photoproduct bands continue to grow steadily and starts to level off after about an hour of
irradiation.
4. The VUV induced photochemistry of the PAH pyrene in a CO ice was also studied.
In this case, PAH ionization is inefficient compared to a VUV irradiated PAH:H2 O
ice and seems to depend on the presence of trace amounts of H2 O.
5. The photochemistry is also tracked as a function of PAH:H2 O ice concentration for
PAH:H2 O ratios ∼1:10 to 1:200. The fraction of neutral molecules converted into
the cation is found to be larger in ices of lower concentration, consistent with the
previous conclusion.
6. The ice photochemistry is monitored for two ice temperatures, 15 and 125 K. PAH
cation formation is important and dominates the first few minutes in the low temperature case. However, ionization is far less efficient in the 125 K ice while the
loss of the neutral parent PAH is equally efficient. It is concluded that reactions
with H2 O photoproducts are the dominant channel at high temperatures.
7. Spectra from 4000 to 500 cm−1 of the non-volatile room temperature residue that
builds up over the coarse of the experiments on each PAH species considered here
are analyzed. These aromatic-rich residues contain OH, CO, and CH2 groups as evidenced by absorptions in the OH stretching region (2.8 µm), CH stretching region
(3.3–3.4 µm) and in the C=O, C–O, and aromatic CC region (5.6–10 µm).
8. The infrared data of the PAH photoproducts are compared to 5–8 µm ISO and
Spitzer spectra of a high- and low-mass protostar. It is shown that PAH:H2 O ice
photoproducts are plausible candidates for the 6.2 to 6.3 µm absorption band associated with interstellar ice and that these contribute to some of the weak structure
between about 6.9 to 8 µm. The band strength of the 6.2 to 6.3 µm band of the
PAH:H2 O ice photoproducts reported here is determined. From this we derive an
upper limit for interstellar abundance of PAHs and PAH:H2 O ice photoproducts
with respect to H2 O ice of 2–3%.
105
Part II
Near-UV/VIS absorption
spectroscopy
107
CHAPTER 5
Real time optical spectroscopy of VUV irradiated
pyrene:H2O interstellar ice1
This chapter describes a near-UV/VIS study of a pyrene:H2 O interstellar ice analogue
at 10 K using optical absorption spectroscopy. A new experimental approach makes it
possible to irradiate the sample with vacuum ultraviolet (VUV) light (7–10.5 eV) while
simultaneously recording spectra in the 240–1000 nm range with sub-second time resolution. Both spectroscopic and kinetic information on VUV processed ices are obtained in
this way. This provides a powerful tool to follow, in situ and in real time, the photophysical and photochemical processes induced by VUV irradiation of a polycyclic aromatic
hydrocarbon containing inter- and circumstellar ice analogue. Results on the VUV photolysis of a prototype sample — strongly diluted pyrene in H2 O ice — are presented.
In addition to the pyrene cation (Py+ ), other products — hydroxypyrene (PyOH), possibly hydroxypyrene cation (PyOH+ ), and pyrene/pyrenolate anion (Py− /PyO− ) — are
observed. It is found that the charge remains localized in the ice, also after the VUV
irradiation is stopped. The astrochemical implications and observational constraints are
discussed.
1 Based on: J. Bouwman, D. M. Paardekooper, H. M. Cuppen, H. Linnartz, L. J. Allamandola, Astrophysical
Journal, 700, 56-62 (2009)
109
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
5.1 Introduction
At present, more than 150 different inter- and circumstellar molecules have been observed
in space. The chemical diversity is striking, and both simple and very complex as well as
stable and transient species have been detected. Among these unambiguously identified
species polycyclic aromatic hydrocarbon molecules (or PAHs) are lacking even though
PAHs are generally thought to be ubiquitous in space [e.g., van Dishoeck 2004]. Strong
infrared emission features at 3.3, 6.2, 7.7, 8.6, and 11.2 µm are common in regions of,
for example, massive star formation and have been explained by PAH emission upon
electronic excitation by vacuum ultraviolet (VUV) radiation. Consequently, PAHs are
expected to play a key role in the heating of neutral gas through the photoelectric effect.
PAHs are also considered as important charge carriers inside dense molecular clouds, and
relevant for molecule formation through ion — molecule interactions [Gillett et al. 1973,
Puget & Leger 1989, Allamandola et al. 1989, Kim et al. 2001, Smith et al. 2007, Draine
et al. 2007]. Nevertheless, the only aromatic species unambiguously identified in space is
benzene, following infrared observations [Cernicharo et al. 2001].
In recent years electronic transitions of PAH cations have been studied in the gas
phase with the goal to link laboratory data to unidentified optical absorption features observed through diffuse interstellar clouds. Following matrix isolation spectroscopic work
[Salama & Allamandola 1991], gas phase optical spectra have been recorded for several
PAH cations [Romanini et al. 1999, Bréchignac & Pino 1999] by combining sensitive
spectroscopic techniques and special plasma expansions [Motylewski et al. 2000, Linnartz 2009]. Such optical spectra have unique features and therefore provide a powerful
tool for identifying PAHs in space. So far, however, no overlap has been found between
laboratory spectra of gaseous PAHs+ and astronomical features.
In dense molecular clouds, most PAHs should quickly condense onto the H2 O- rich
icy grain mantles, quenching the IR emission process. Here, they will participate in ice
grain chemistry. More than 25 years of dedicated studies, mainly in the infrared, have
proven that a direct comparison between laboratory and astronomical ice spectra paints
an accurate picture of the composition and the presence of inter- and circumstellar ices,
even though solid-state features are rather broad. The spectral features (band position,
band width (FWHM) and the intensity ratio of fundamental vibrations) depend strongly
on mixing ratio and ice matrix conditions and this provides a sensitive analytical tool to
identify ice compositions in space [e.g., Boogert et al. 2008, Öberg et al. 2008].
In the past, several experiments have been reported in which the formation of new
molecules was proven upon VUV irradiation of astronomical ice mixtures, typically under high vacuum conditions [e.g., Mendoza-Gomez et al. 1995, Bernstein et al. 1999,
Gudipati & Allamandola 2003, Ruiterkamp et al. 2005, Peeters et al. 2005, Elsila et al.
2006]. Many of these studies were not in situ, i.e., reactants were determined after warm
up of the ice, and although VUV-induced photochemistry at low temperatures is expected
to take place, it is not possible to fully exclude that at least some of the observed reactants
may have been formed during the warm-up stage. More recently, in situ studies have
become possible using ultrahigh vacuum setups in which ices are grown with monolayer
precision and reactions are monitored using reflection absorption infrared spectroscopy
110
5.2 Experimental
and temperature programmed desorption. Recent results show that H-atom bombardment
of CO and O2 ice results in the efficient formation of H2 CO/CH3 OH and H2 O2 /H2 O, respectively [Watanabe & Kouchi 2002, Ioppolo et al. 2008, Miyauchi et al. 2008]. Overall,
such studies are still restricted to the formation of rather small species, with ethanol as
the most complex molecule investigated in this way [Bisschop et al. 2007b]. Moreover,
the reactants and products should not have overlapping bands in order to track them separately.
In the present work, a new approach is presented that extends our previous FTIR work
on interstellar ice analogues to the UV/VIS. With the new setup it is possible to record,
in situ, the VUV photochemistry of PAHs and PAH derivatives in water ice at 10 K in
real time. In the next section, this new approach is discussed in detail. The first results for the PAH pyrene (C16 H10 , or Py) and its photoproducts, pyrene cation (C16 H+10 ,
or Py+ ), hydroxypyrene (PyOH), hydroxypyrene cation (PyOH+ ), and pyrene/pyrenolate
anion (Py− /PyO− ) in H2 O ice are presented in §5.3. Finally, the astrophysical relevance of
this work is discussed in §5.5. The latter is twofold. First, time-dependent results of VUV
irradiated ice provide general insight into possible reaction pathways upon photoprocessing of PAH containing water ice. Second, the results provide a spectroscopic alternative
to search for PAH and PAH-related optical features in the inter- and circumstellar medium
(ISM/CSM) through electronic solid-state absorptions.
5.2 Experimental
A schematic of the experimental setup is shown in Fig. 5.1. The experiment consists of
three units: a vacuum chamber in which the ice is grown, a special VUV irradiation source
that is used for the photo-processing of the ice and a source emitting broadband light that
is focused into the ice and subsequently detected using a monochromator equipped with
a sensitive CCD camera.
The vacuum chamber consists of an ISO-160 6-cross piece. A 300 l s−1 turbomolecular pump, backed by a 10 m3 hr−1 double stage rotary pump, is used to evacuate the chamber and to guarantee an operating pressure of ∼10−7 mbar. A catalytic trap is mounted on
the pre-vacuum pump to prevent pump oil from entering the vacuum chamber.
The top flange of the cross piece holds a differentially pumped rotary flange on which
is mounted a closed cycle helium refrigerator equipped with a cold finger. A MgF2 sample
window with a diameter of 14 mm, clamped into an oxygen free copper holder between
indium gaskets, is mounted on the cold finger and centered on the optical axis of the setup.
This allows for rotation of the sample window through 360◦ under vacuum. The sample
window can be cooled down to 10 K and a thermocouple (Chromel-Au/Fe (0.07%)) and
a temperature controller guarantee accurate temperature settings with 0.1 K precision.
The Py:H2 O sample is prepared by vapor depositing pyrene from a solid sample
(Aldrich 99%) heated to 40◦ C, together with milli-Q water vapor from a liquid sample.
The entire inlet system is maintained at ≥40◦ C during deposition and comprises gas bulbs
containing the sample material and tubing for directed deposition, approximately 15 mm
from the sample window. The flow rate of the sample material is set by a high precision
111
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
M2
U
L3
P1
Hydrogen
H2−lamp
Xe− or Fiber−lamp L1
Temperature
Controller
T1
ID3
W3
W1
ID2
T1
Computer
Sample
Pump
ID1
HeNe
CCD
PD
M1
MW−
Power
Supply
W2
SW
VC
Pump
L2
Spectr.
BS
P1
Computer
Figure 5.1 A schematic drawing of the experimental setup. BS: beam splitter; IDX: iris
diaphragm X; LX: lens X; MX: mirror X; PD: photodiode to monitor interference fringes;
PI: pressure indicator; SW: MgF2 cold sample window; TI: temperature indicator; U: voltage meter; VC: Vacuum Chamber and WX: MgF2 window X. The light paths are indicated
by arrows, the data wiring is indicated in red lines, and the hydrogen flow is indicated by
blue lines (see online color version). The inset shows a three-dimensional drawing of the
experimental setup.
dosing valve. Condensation inside the tube is prohibited by additional resistive heating
and the temperature settings are monitored by K-type thermocouples. The resulting ice
film thickness is accurately measured by recording the number of interference fringe maxima (m) of a HeNe laser (λ =632.8 nm) which strikes the sample window at an angle of
θ=45◦ . To monitor film growth and thickness, the intensity of the reflected laser light is
measured with a sensitive photodiode. The ice thickness is subsequently determined by:
d=
mλHeNe
,
2nice cos θ
(5.1)
with the refractive index of the predominantly H2 O ice being nice ≈ 1.3 [Hudgins et al.
1993]. This is illustrated in Fig. 5.2 where both interference fringes produced during
ice deposition and the simultaneous growth of the integrated neutral Py absorbance band
112
5.2 Experimental
are shown. The final ice thickness amounts to 1.7 µm and is reproducible to within 5% or
better. Simultaneously, the number of pyrene molecules in the ice sample (N) is monitored
by measuring the integrated absorbance of its strongest transition (S2 ← S0 ) (see Fig. 5.2).
The number of pyrene molecules per cm2 can be calculated via [Kjaergaard et al. 2000,
Hudgins et al. 1993]:
R ν2
τ dν
ν1
N=
,
(5.2)
8.88 × 10−13 f
where f = 0.33 is the known oscillator strength of the S2 ← S0 transition of pyrene [Bito
et al. 2000, Wang et al. 2003]. The resulting column density of pyrene molecules amounts
to about 4×1014 cm−2 . For a typical sample with a thickness of 1.7 µm, the column density
of H2 O molecules amounts to 4×1018 cm−2 , using the value for the density of amorphous
ice (ρ=0.94 g/cm3 ; Sceats, M. G. and Rice, S., A. [1982]). Thus, the sample is calculated
to consist of a 1:10,000 pyrene:H2 O mixture. This mixing ratio can be roughly varied by
changing the H2 O flowrate or the pyrene sample temperature. The HeNe beam used here
for monitoring the ice growth process also traces other elements along the optical path
and is used to align all components.
The vacuum UV radiation from a special microwave (MW) powered hydrogen discharge lamp is used to simulate the interstellar radiation field [Muñoz Caro et al. 2002].
100
-0.20
-0.21
60
-0.22
40
20
-0.23
Photodiode voltage / V
Integrated absorbance / cm
-1
80
error
0
-0.24
0
2
4
6
8
10
12
14
16
Time / minutes
Figure 5.2 A plot showing both the growth of the integrated pyrene absorption band
(squares and left axis) and the interference fringes measured by the photodiode (right
axis). At t = 14 minutes, the deposition is stopped and the interference pattern diminishes. The error bar shown in the right lower corner applies to the pyrene-integrated
absorbance.
113
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
The lamp consists of a flow tube clamped in a McCarroll cavity [McCarroll 1970] and
emits mainly Ly-α radiation around 121.6 nm and, with less intensity, a band centered
around 160 nm. The cavity is excited by a regular MW power supply (100 W, 2450 MHz).
The H2 pressure in the lamp is maintained at 0.4 mbar during operation (Praxair 5.0 H2 ).
This results in a VUV photon flux of ∼1015 photons cm−2 s−1 . The lamp is centered onto
the front flange and the VUV radiation enters the setup towards the ice sample through a
MgF2 window that also serves as a vacuum seal. A shutter is used to block the VUV until
the moment that the ice processing should start. Besides eliminating the need to switch
the H2 lamp on and off during the course of an experiment, this allows the lamp to stabilize before irradiation starts. This is important when tracking photochemical behavior
during extended periods of photolysis.
A 300 W ozone-free Xe-arc lamp serves as a broad band white light source to measure
the spectral ice features in direct absorption. The lamp has a spectral energy distribution
that covers the full detector range (200 nm < λ < 2400 nm). Alternatively, the light from
a halogen fiber lamp can be used when no UV coverage is desired. An optical system consisting of lenses and diaphragms is used to guide the light beam through a MgF2 window
along the optical axis — coinciding with the pre-aligned HeNe beam — and crossing the
ice sample at a 45◦ angle. Light that is not absorbed exits the vacuum chamber through
a second MgF2 window after which it is focused onto the entrance slit of an ANDOR
Shamrock spectrometer. The spectrometer is equipped with two interchangeable turrets
which holds four gratings in total (2400, 1200, 600 and 150 lines mm−1 ), allowing for a
trade-off between wavelength coverage and spectral resolution, depending on the experimental needs. Since typical ice absorption bands exhibit a FWHM of 4–20 nm, most of
the experiments are performed using the 600 lines mm−1 grating, resulting in an accessible
wavelength range of ∼140 nm.
The light is dispersed onto a very sensitive 1024×256 pixel CCD camera with 16bit digitization. The resulting signal is read out in the vertical binning mode by a data
acquisition computer. Spectra are taken in absorbance mode (τ=-ln(I/I0 )) with respect to
a reference spectrum (I0 ) taken directly after depositing the sample. Recording a single
spectrum typically takes about 5 ms and spectra are generally co-added to improve the
signal-to-noise ratio (S/N). In a typical experiment more than 1000 individual spectra are
recorded and are reduced using LabView routines. Data reduction consists of local linear
baseline corrections, multiple Gaussian fitting of absorption profiles and absorption band
integration.
5.3 Spectroscopic assignment
Figure 5.3 shows the 310 to 500 nm spectrum of a Py:H2 O ice at 10 K after 1200 s of
in situ photolysis in absorbance mode. The spectrum is baseline subtracted and given
in optical depth (OD). Since the spectrum recorded before VUV irradiation is taken as a
reference (I0 ), bands with positive OD values arise from species produced by photolysis
while the carriers of negative OD bands decrease in density. It is noteworthy that S/N
ratios are good even though the processes are studied in a very dilute mixture (Py:H2 O
114
5.3 Spectroscopic assignment
∼1:10,000). Previous work was on more concentrated samples (PAH:H2 O ∼1:500, Gudipati & Allamandola [2003]; PAH:H2 O ∼1:800 to 1:3200, Bernstein et al. [1999]).
A Gaussian fit to all of the features visible in the spectrum is indicated as well. Clearly,
a number of new species are produced at the expense of neutral pyrene. The peak positions, FWHM and assignments of all the bands in Fig. 5.3 are summarized in Table 5.1,
along with comparisons of earlier results found in other molecular environments. The
assignments given in Fig. 5.3 were made as follows. Based on previous studies of pyrene
in rare gas matrices [Vala et al. 1994, Halasinski et al. 2005], the strong, negative band
peaking at 334 nm is readily assigned to the 1 B2u ←1 Ag electronic transition of neutral
pyrene (S 2 ←S 0 ).
Similarly, the positive bands near 363, 446, and 490 nm are assigned as the strongest
members of the pyrene cation 2 B1u ←2 B3g , 2 Au ←2 B3g and 2 B1u ←2 B3g vibronic transitions,
respectively [Vala et al. 1994, Hirata et al. 1999]. Table 5.1 shows that the bandwidth
(FWHM) of the Py+ bands is broader in the solid H2 O than in rare gas matrices, in accordance with the stronger interactions within the H2 O matrix network. Similarly, larger
shifts in peak position may be expected.
In addition to the Py+ bands, other new bands appear near 345, 367, 405, and 453 nm.
We ascribe these to hydroxypyrene (PyOH), hydroxypyrene cation (PyOH+ ), and pyrene/
pyrenolate anion (Py− /PyO− ) based on the work of Milosavljevic & Thomas [2002] who
reported the spectra of 1-hydroxypyrene and its daughter products in various media. The
suggestion of anion production in these ices is noteworthy in view of the astronomical
detection of negative ions both in the solid state [van Broekhuizen et al. 2005] and in the
gas phase [e.g., Agúndez et al. 2008].
The appearance of clear bands due to PyOH, PyOH+ , and Py− /PyO− after VUV radiation at 10 K is somewhat surprising. Previous optical studies of the VUV photolysis
of three different PAHs in H2 O ice indicated that conversion of the parent neutral PAH
to the cation was the major, and apparently only, photolytic step and that the cation remained stabilized in the ice to remarkably high temperatures (∼100 K) for long periods
[e.g., Gudipati & Allamandola 2006b]. Gudipati [2004] further showed that, in the case
of naphthalene (Nap), subsequent reactions between Nap+ and the H2 O matrix did indeed
produce NapOH, but only during warm-up (∼100 K). Using mass spectroscopy, Bernstein et al. [1999] showed that PAH oxides and hydroxides were part of the residues left
after VUV photolyzed PAH:H2 O ices were warmed under vacuum. Furthermore, upon
prolonged exposure, the VUV transmittance of the MgF2 hydrogen lamp window drops
and with this the PyOH and PyOH+ production also decreases relative to the production
of Py+ . This suggests that photolytic processes within the pyrene containing water ice
change, perhaps because water dissociation becomes less effective with reduced hard UV
flux while direct Py ionization still readily occurs with near UV photons [Gudipati & Allamandola 2004]. The influence of temperature and UV spectral energy distribution is
currently under investigation.
115
31
30
29
28
1
B
2u
26
25
-1
24
23
22
21
20
+
Py
Py
0.02
27
2
1
B
A
1u
g
2
B
3g
Production
Optical Depth (-)
0.01
0.00
2
-0.01
-
+
PyOH
PyOH
A
-
2
2
+
PyOH
B
u
3g
2
B
1u
B
3g
+
Py
+
Py
Py / PyO
-0.02
Pyrene: C
16
H
10
-0.03
-0.04
Depletion
-0.05
320
340
360
380
400
420
440
460
480
500
Wavelength / nm
Figure 5.3 A baseline corrected spectrum obtained after 1200 s of VUV irradiation of a pyrene:H2 O ice at 10 K. The absorption
features in the plot are assigned and fitted with Gaussian profiles. A negative OD indicates species destruction, a positive OD species
formation. The left spectrum (up to 280 nm) is scaled down by 30% to facilitate comparison.
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
116
3
Wavenumber / 10 cm
32
Table 5.1 Vibronic bands of pyrene and its photoproducts in a H2 O ice compared with rare gas matrix literature values.
Species
Py
State
1B
2u
Py+
2B
1u
Py+
2A
u
Py+
2B
1u
λH2 O (nm)
334.0
329.2
319.2
363.2
354.0
344.9
445.6
435.5
423.0
413.8
490.1
FWHM (nm)
4.4
3.2
6.5
3.6
6.5
6.2
6.6
10.2
12.2
5.3
10.0
λlit. (nm)
323.3a
319.1a
309.4a
362.6b
355.1b
443.8b
433.2b
422.9b
412.1b
486.9b
FWHMlit. (nm)
n.a.
n.a.
n.a.
2.2
2.1
4.5
4.3
4.1
3.9
5.5
∆FWHM
1.4
4.4
2.1
5.9
8.1
1.4
4.5
-
-
117
5.3 Spectroscopic assignment
PyOH
344.9
5.8
340c
Py− /PyO−
405.2
7.3
410c
PyOH+
366.8
3.0
PyOH+
452.9
18.2
465c
a Values measured in a Ne matrix taken from Halasinski et al. [2005]
b Values measured in an Ar matrix taken from Vala et al. [1994]
c Values measured in H O and 2-chlorobutane taken from Milosavljevic & Thomas [2002]
2
λH2 O − λrg
10.7
10.1
9.8
0.6
-1.1
1.8
2.3
0.1
1.7
3.2
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
5.4 Chemical evolution of the ice
To further investigate the spectroscopy and the photochemistry of VUV irradiated H2 Orich ices that contain PAHs, time-dependent optical studies were performed. Figure 5.4
shows the integrated OD behavior of the Py, Py+ (for two bands), and PyOH absorptions
as function of photolysis time. During the first 130 s of VUV irradiation, the Py decay is
clearly correlated with Py+ growth. This allows us to determine relative band strengths of
the two species by investigation of the short timescale correlation. The sub-second time
response of the present setup is a prerequisite for this to work. We derive a band strength
of 2.9 × 10−13 cm molecule−1 for the 2 A2u ← 2 B3g Py+ transition in H2 O ice using Equation (5.2). As other chemical processes become important, the correlation disappears.
The loss of Py slows significantly while the Py+ starts a slow decline after the maximum
is reached. The PyOH signal continues to grow slowly but steadily throughout the photolysis process and is most likely formed by Py/Py+ reacting with photoproducts of H2 O. Its
formation is consistent with the recent outcome of a quantitative VUV photodesorption
study of H2 O ice under ultrahigh vacuum conditions, where a H+OH photodissociation
channel was reported [Öberg et al. 2009d]. A reaction network connecting all of these
species is presented in Fig. 5.5.
VUV fluence / 10
0
2
4
18
photons
6
8
10
50
Integrated O.D. / cm
-1
Py
40
PyOH
30
20
+
Py
10
0
+
Py
0
20
40
60
80
100
120
140
160
Photolysis time / minutes
Figure 5.4 The behavior of Py, Py+ (two transitions), and PyOH as a function of VUV
photolysis time/VUV fluence.
In addition to Py+ formation and reactivity during irradiation, we have also studied
its stability within the ice when photolysis is stopped. Figure 5.6 plots the normalized
integrated O.D. of the 445.6 nm Py+ band as a function of time after VUV radiation is
118
5.4 Chemical evolution of the ice
e-
Py-
e-
Py
OH
PyO-
e-
h
PyO
Py+
h
OH
e-
PyOH
PyOH+
h
Figure 5.5 Possible reactions upon photolysis of Py:H2 O ice as derived from Fig. 5.3.
stopped. The figure spans 50 hr and shows that although the total Py+ signal drops, 60%
remains trapped in the ice after 2 days. The small wiggle at the 0.05 level is due to
baseline variations. There are clearly two decay channels, one ‘fast’ and one ‘slow’. The
following expression is used to fit the experimental data:
y = A1 exp(−t/τ1 ) + A2 exp(−t/τ2 )
(5.3)
with A1 = 0.70, τ1 = 351.3 hr, A2 = 0.22 and τ2 = 1.7 hr. This produces the red curve in
Fig. 5.6. The processes responsible for these two decay rates are not yet clear. The Py+
decay may be governed by recombination with trapped electrons. In considering these
results, it is important to keep in mind that the ice processes described here are recorded
for one temperature (10 K) and will most likely depend on temperature.
Normalized integrated O.D.
1.0
0.9
0.8
0.7
0.6
0.5
error
0.4
0.3
0
5
10
15
20
25
30
35
40
45
Time / hours
Figure 5.6 Normalized pyrene cation integrated O.D. as a function of time plotted together
with a double exponential fit using Equation (5.3) (red curve). The wiggles superposed
on the signal are caused by baseline effects and fall within the error of 5% as indicated by
the error bar.
119
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
5.5 Astrophysical implications
Water is by far the dominant component of interstellar ices. Since PAHs are considered
to be widespread throughout the ISM, they are likely to be frozen out wherever H2 O-rich
ices are present. The photochemical kinetics observed here and the new spectroscopic
information make two astrophysical points.
Astrochemically, this work shows that the effective photolytic production of PAH ions
in PAH containing ice upon VUV irradiation should not be neglected a priori when modelling interstellar ice chemistry. The behavior of the various species that is shown in
Fig. 5.4 suggests that a new set of solid-state reactions appears when irradiating PAH containing water ice. The present study is on a rather isolated ice system — typical for this
type of laboratory study — comprising Py and H2 O. In a more realistic interstellar sample, containing other constituents, such as CO, CO2 or NH3 , and other PAHs, chemical
pathways will become more complicated, but since water is the most dominant component in these extraterrestrial ices, we expect that the trends observed here will generally
apply. Another important point to note is that reactions involving ions are not included in
any of the current astrochemical grain chemistry networks. The present study shows that
positive ions can reside in the ice mantle for a substantial time. This is particularly interesting since an astronomical dust grain is a truly isolated system, whereas the laboratory
analogue is grown on the tip of a cold finger.
Observationally, the spectroscopic results provide an alternative route to search for
PAH features in space. Astrophysical searches to identify PAHs and PAH cations in the
ISM/CSM have focused on vibrational and electronic transitions in the gas phase as well
as solid-state PAH features in the infrared. As stated previously, these searches have been
largely unsuccessful. The infrared work suffers from spectral congestion and spectral
overlap of vibrational modes. This prohibits an unambiguous identification of an emission
feature to a specific carrier. In the UV/VIS this is partly overcome as electronic excitations
are unique for different PAHs. However, electronic spectra of gas phase PAHs showed no
overlap with absorption features recorded through diffuse interstellar clouds, presumably
of too low column densities [Romanini et al. 1999, Bréchignac & Pino 1999]. A different
situation applies in the solid state; PAHs are refractory material and accumulate in time
onto cold grains that offer a reservoir. Therefore, the specific embedding of PAHs in water
ice as presented here provides an alternative starting point for an astronomical search.
However, one has to realize that this idea has both pros and cons.
Table 5.1 shows that under the present experimental conditions the bands are rather
broad. The FWHM of the Py+ absorption feature at 445.6 nm is 66 Å. From an observational point of view this has the advantage that profiles can be studied at medium resolution, but also goes with the challenge to correct very accurately for possible background
signals and in the end spectral overlaps may still exist, e.g., with silicate or carbonaceous
features. Nevertheless, overlapping broad bands can contribute to the very broad structure (VBS) superposed on the interstellar extinction curve [Hayes et al. 1973, van Breda
& Whittet 1981, Krelowski et al. 1986] and simply may have been overlooked in the past.
As stated before, the electronic excitation energy is unique. This is good for selectivity, but bad for sensitivity, since spectral features of different PAHs do not add up (as in
120
5.6 Conclusion
the infrared) and consequently an individual optical band strength is, in principle, directly
determined by the actual abundance of one specific PAH in space. It is difficult, however,
to predict the percentage of the total interstellar PAH population that might exist in the
form of one specific PAH, e.g., pyrene. Nevertheless, electronic PAH transitions are typically 2–3 orders of magnitude stronger than their IR bands and from this point of view,
we expect that PAH bands may be observable in the visible even though IR bands are
barely discernible on the strong H2 O bands [Brooke et al. 1999]. This is also reflected by
the very good S/N ratios as visible in Fig. 5.2 for a very diluted mixture. We expect PAHs
to be sufficiently abundant in ices in regions of molecular clouds with Av ≥8 to permit
detection in the optical. Infrared ice bands have been detected along such lines of sight
and the visible extinction is low enough to permit UV from the interstellar radiation field
to process these ices.
We have estimated the expected Py and Py+ absorption band strength towards an
example source, MWC297. This is an early-type B1.5V star with a well-characterized
stellar spectrum, a B magnitude of 14.34 and a 3.0 µm ice band with τ = 0.04. Using the
high sensitivity of an 8 m class telescope (VLT) it is feasible to obtain, within a couple of
hours, S/N ≥1000 spectra of such a highly extincted source in the wavelength range under
investigation. To estimate the expected OD of a pyrene or pyrene cation ice absorption
we use the standard relation:
NH /E(B − V) = 5.8 × 1021 atoms cm−2 mag−1
(5.4)
from Bohlin et al. [1978]. With E(B−V) = 2.67 towards MWC297 [Drew et al. 1997] this
results in NH = 1.6×1022 cm−2 . Taking the total PAH abundance in clouds with respect to
nH to be ∼3×10−7 , this results in a total PAH column density of 9.6×1015 cm−2 . Assuming
that 1% of the total PAHs in space is in the form of pyrene frozen out on grains and of this
fraction up to 10% is in its singly ionized state (Py+ ), the total column density of Py+ or
Py toward MWC297 is estimated to range between 9.6×1012 and 8.64×1013 species/cm2 .
The OD is defined as
NA
,
(5.5)
τ=
∆ν
with N the column density of the absorbers, A the integrated band strength and ∆ν the
FWHM. For a typical strong allowed vibronic transition, such as the 1 B2u ←1 Ag Py and
2
Au ←2 B3g Py+ transitions, we take APy = 2.9 × 10−13 cm molecule−1 and APy+ = 2.9 ×
10−13 cm radical−1 in ice with a FWHM=400 cm−1 and 300 cm−1 , respectively. This
yields ODs of 0.01 ≤ τ ≤ 0.06 for Py and Py+ , a range similar to that observed for ice
bands.
5.6 Conclusion
This work presents the first results of a spectroscopic and photochemical study of pyrene
in water ice upon VUV irradiation under astronomical conditions. Since the spectra are
recorded in real time, it is possible to derive photochemical characteristics and to monitor
a rich ion-mediated chemistry in the solid state. Such processes are yet to be considered in
121
5 Optical spectroscopy of VUV irradiated pyrene:H2 O ice
astrochemical models. Additionally, it is shown that the pyrene cations formed within the
H2 O ice by VUV irradiation remain trapped in the ice for an extended period. Successive
heating of the ice makes these ions available to diffusing species and hence should be
considered in solid-state astrochemical processes.
The new laboratory approach presented here offers a general way to provide astronomically relevant PAH solid-state spectra. Specifically, the spectra discussed here provide
an alternative way to search for pyrene features in the ISM/CSM. The derived numbers
show the potential of this method, but one has to realize, as pointed out before, that these
numbers incorporate our limited knowledge on the actual PAH quantities in space. For
different PAHs, with different abundances and different absorption strengths other numbers, both less and more favorable, may be expected. Furthermore, it is possible — in
view of the rather effective way in which charged species form and stay in the ice — that
the actual abundance of ions may be higher. The results presented here are new and aim
at a further characterization of the chemical role of PAHs and PAH derivatives in space.
122
CHAPTER 6
Photochemistry of the PAH pyrene in water ice:
the case for ion-mediated solid-state
astrochemistry1
Icy dust grains play an important role in the formation of complex inter- and circumstellar molecules. Laboratory studies have mainly focused on the physical interactions and
chemical pathways in ices containing rather simple molecules, such as H2 O, CO, CO2 ,
CH4 , and CH3 OH. Observational studies show that polycyclic aromatic hydrocarbons
(PAHs) are also abundantly present in the ISM in the gas phase. It is likely that these
non-volatile species also freeze-out onto dust grains and participate in the astrochemical
solid-state network, but additional experimental PAH ice studies are largely lacking. The
study presented here focuses on a rather small PAH, pyrene (C16 H10 ), and aims to understand and quantify photochemical reactions of PAHs in interstellar ices upon vacuum
ultraviolet (VUV) irradiation as a function of astronomically relevant parameters. NearUV/VIS spectroscopy is used to track the in situ VUV driven photochemistry of pyrene
containing ices at temperatures ranging from 10 to 125 K. The main photoproducts of
VUV photolyzed pyrene ices are spectroscopically identified and their band positions are
listed for two host ices, H2 O and CO. Pyrene ionization is found to be most efficient in
H2 O ices at low temperatures. The reaction products, triplet pyrene and the 1-hydro-1pyrenyl radical are most efficiently formed in higher temperature water ices and in low
temperature CO ice. Formation routes and band strength information of the identified
species are discussed. Additionally, the oscillator strengths of Py, Py·+ , and PyH· are
derived and a quantitative kinetic analysis is performed by fitting a chemical reaction
network to the experimental data. The results are placed in an astrophysical context by
determining the importance of PAH ionization in a molecular cloud. The photoprocessing
of a sample PAH in ice described in this chapter indicates that PAH photoprocessing in
the solid state should also be taken into account in astrochemical models.
1 Based on: J. Bouwman, H. M. Cuppen, A. Bakker, L. J. Allamandola, and H. Linnartz Astronomy and
Astrophysics, 511, A33+ (2010)
123
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
6.1 Introduction
Strong infrared emission attributed to polycyclic aromatic hydrocarbons (PAHs) is characteristic of many galactic and extragalactic objects [Smith et al. 2007, Draine et al.
2007, Draine & Li 2007, Tielens 2008]. While this emission generally originates in
optically thin, diffuse regions, PAHs should also be common throughout the dense interstellar medium. There, as with most other interstellar species in molecular clouds,
PAHs condense out of the gas onto cold icy grain mantles, where they are expected to
influence or participate in the chemistry and physics of the ice. While laboratory studies
of interstellar ice analogs have shown that complex organic molecules are produced upon
extended vacuum ultraviolet (VUV) photolysis [e.g., Briggs et al. 1992, Bernstein et al.
1995], the photoinduced processes occurring during the irradiation of PAH containing interstellar ice analogues have not yet been studied in detail. In optical, in situ studies of
the photochemistry of naphthalene, 4-methylpyrene, and quatterylene containing water
ice at 20 K, Gudipati & Allamandola [2003, 2006a,b] and Gudipati [2004] showed that
these PAHs are readily ionized and stabilized within the ice, suggesting that trapped ions
may play important, but overlooked roles in cosmic ice processes. Beyond this, there is
little information about the VUV induced, in situ photochemistry and photophysics of
PAH-containing water-rich ices.
Here, we describe a detailed study of the VUV-induced photochemistry that takes
place within pyrene (Py or C16 H10 ) containing water ices (Py:H2 O=1:10,000–1:5,000).
The present study is an extension of Chapter 5 in which the focus was on the new experimental setup and where the use of PAH ice spectra was discussed to search for solid-state
features of PAHs in space. In this work, the focus is on a detailed characterization of
the chemical processes taking place upon VUV irradiation, particularly as a function of
ice temperature ranging from 25 to 125 K. Additionally, measurements on Py:CO ices
at 10 K were performed to elucidate the role of water in the reaction schemes and to
clarify the formation routes of identified species. A similar study of three small PAHs
is now underway to understand the general principles of PAH/ice photochemistry. This
is part of an overall experimental program at the Sackler Laboratory for Astrophysics to
study the fundamental processes of inter- and circumstellar ice analogues such as thermal
[Acharyya et al. 2007] and photodesorption [Öberg et al. 2007b, 2009d], hydogenation
reactions [Fuchs et al. 2009, Ioppolo et al. 2008], photochemistry [Öberg et al. 2009c],
and physical interactions in interstellar ice analogues [Öberg et al. 2007a, 2009b, and
Chapter 2 of this thesis].
The chapter is organized as follows. The experimental technique is summarized in
§6.2. Paragraph 6.3 describes the Py:H2 O and Py:CO ice photochemistry, the resulting products and their formation routes. The temperature-dependent photochemistry and
derived reaction dynamics are described in §6.4 and astrochemical implications are discussed in §6.5. The main conclusions are summarized in §6.6.
124
6.2 Experimental technique
6.2 Experimental technique
We use the new apparatus as described in Chapter 5 which follows the photochemistry in
kinetic mode during VUV irradiation by measuring the near-UV-visible absorption spectra of an ice, providing ‘real-time’ tracking of the reactants and photoproducts. Dilute
Py:H2 O ice samples (∼1:10,000–∼1:5,000) and a Py:CO ice sample of comparable concentration are prepared by depositing the vapor from a pyrene sample heated to 40◦ C
together with H2 O vapor or CO gas onto a cold MgF2 window. The window is cooled
to 10 K in the case of CO deposition or 25 K in the case of H2 O deposition. The sample window is cooled by a closed cycle He refrigerator. Pyrene (Aldrich, 99%) and CO
(Praxair 99.999%) are used as commercially available. Vapor from water, filtered through
a milli-Q purification system and purified further by three freeze-pump-thaw cycles, is
used. The sample window is mounted in a high-vacuum chamber (P ≈ 10−7 mbar).
The ice growth rate and thickness are determined with a HeNe laser by monitoring the
thin-film interference fringes generated during deposition. Simultaneously, the amount of
pyrene is tracked by measuring the integrated strength of the S2 ← S0 neutral Py transition
at 334 nm. Deposition is typically stopped when the optical depth (OD) of Py approaches
∼0.15.
The ice samples are photolyzed with the 121.6 nm Lyα (10.6 eV) and the 160 nm
molecular hydrogen emission bands (centered around 7.8 eV) generated by a microwave
powered discharge in a flowing H2 gas with a VUV flux of ∼ 1015 photons s−1 [Muñoz
Caro et al. 2002]. This results in a photon flux of ∼ 1014 photons·cm−2 s−1 at the sample
surface [Öberg et al. 2009d].
Absorption spectra of VUV-photolyzed Py-containing ices are measured with a Xearc lamp serving as a white light source. Lenses and diaphragms direct the light through
the ice sample along the optical axis determined by the HeNe laser beam after which
it is focused onto the entrance slit of a 0.3 m spectrometer. A 150 lines mm−1 grating,
blazed at 300 nm, disperses the light onto a sensitive 1024×256 pixel CCD camera with
16 bit digitization. The camera is read out in vertical binning mode by a data acquisition
computer that converts the data to absorbance spectra (OD = − ln(I/I0 )). This configuration spans the 270 to 830 nm spectral range, which permits simultaneously monitoring
of the behavior of the neutral Py parent molecule and photoproduct bands without any
adjustment of the elements along the optical path. This is critical to obtaining reliable and
reproducible baselines in measuring the optical spectra of ices. The spectral resolution is
of the order of 0.9 nm, which is more than sufficient to record broad solid-state absorption
features.
The measurements described here were performed on various H2 O:Py ice samples at
25, 50, 75, 100, and 125 K. The CO ice experiments were carried out at 10 K to avoid
matrix sublimation at higher temperatures. The sample temperature is maintained using
a resistive heater with an accuracy of ±2 K. The measured spectra are converted into
units of optical depth by using the spectrum of the freshly deposited, unphotolyzed ice
at the appropriate temperature as a reference spectrum (I0 ). Recording a single spectrum
typically takes about 5 ms, and 229 spectra are generally co-added to improve the S/N of
a spectrum, producing one single spectrum every 10 seconds.
125
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
The optical configuration of the apparatus is such that spectra are recorded simultaneously with photolysis. Thus, the short spectral recording time permits us to monitor
photoinduced changes on a roughly 10 s time scale. Figure 6.1 shows the 290 to 490 nm
spectrum of a Py:H2 O ice at 125 K after 900 s of in situ VUV photolysis. Because the
spectrum recorded before VUV irradiation is taken as a reference (I0 ), bands with positive
OD values originate from species that are produced by photolysis, while the bands with
negative OD correspond to the neutral pyrene that is lost upon photolysis. Comparing the
Py and photoproduct absorption bands with the narrow Hβ lamp line at 486.1 nm shows
that the instrumental resolution indeed far exceeds the ice band widths. The absolute
wavelength calibration is accurate to within ±0.5 nm.
More than 1400 individual spectra are recorded and are reduced in a typical 4 hr experiment. Spectra are individually baseline corrected by fitting a second order polynomial
through data points where no absorption occurs and subsequently subtracting the fit from
the measured spectrum. Integrated absorbances of absorption features are calculated numerically for all spectra. These are corrected for the contributions of atomic hydrogen
lines originating in the H2 discharge lamp. The data reduction software also allows us to
plot correlation diagrams between integrated absorbances of different absorption features.
All data handling and reduction is performed with LabView routines.
Integrated band areas are used, in conjunction with oscillator strengths ( f ), to derive molecular abundances. The oscillator strength is converted to integrated absorbance
(cm molecule−1 ) using the conversion factor 8.88×10−13 [Kjaergaard et al. 2000]. The
number of molecules per cm2 (N) is given by
R ν2
τ dν
ν1
,
(6.1)
N=
8.88 × 10−13 f
where τ is the optical depth and ν is the frequency in cm−1 .
6.3 Band assignments and band strength analysis
The typical photolysis duration of about 4 hours is the time required for nearly complete
loss of the neutral pyrene vibrational progression at 334.0, 329.2, and 319.2 nm. Irradiating the sample ices with VUV light produces a set of new absorption bands in the
spectra, indicating active photochemistry. The band positions, FWHM, and assignments
of the bands in the Py:H2 O ice at 25 K are listed in Table 6.1. The bands appearing in the
Py:CO ice at 10 K are similar to those in the Py:H2 O ice at 25 K, although, with slightly
altered band positions and FWHM and with very different relative intensities (see also
Table 6.1). Figure 6.1 presents a spectrum from the 125 K Py:H2 O series. This figure
illustrates production of the pyrene radical cation (Py·+ ), triplet pyrene (3 Py), 1-hydro-1pyrenyl radical (PyH· ), and a broad underlying ‘residue’ feature upon VUV irradiation.
Additionally, a progression of distinct absorption features is found in the Py:CO experiment, which indicates the formation of the (reactive intermediate) HCO· radical. The
identifications of these species and their oscillator strengths are discussed below.
126
3
Wavenumber / 10 cm
34
32
30
28
-1
26
.+
24
PyH
Py
.
3
22
.+
Py
Py
Py
H
Py photoproducts
-0.05
-0.10
300
320
340
360
380
400
420
440
460
480
Wavelength / nm
127
Figure 6.1 The spectrum of a dilute pyrene:H2 O ice after 900 s of VUV irradiation at 125 K. The inset shows a blow-up of the
pyrene photoproduct bands. Band assignments are discussed in §6.3. Note the broad feature ranging from about 350 to 470 nm
which is indicated by a Gaussian fit. This is attributed to overlapping bands from individual pyrene photoproducts. Bands with
negative optical depth indicate species destruction, those with positive optical depth show species formation. The blue bands are
Gaussian profiles which co-add to the overall fit shown in red. Note the instrumental resolution indicated by the profile of the Hβ
line at 486.1 nm. (This figure is available in color in electronic form)
6.3 Band assignments and band strength analysis
Optical Depth (-)
0.00
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
Table 6.1 Band positions (λc ) and FWHM in nm for pure pyrene ice at 10 K, pyrene in
H2 O ice at 25 K, pyrene in CO ice at 10 K, and photoproduct bands for the Py:H2 O and
Py:CO UV processed ices.
Species
Py 1 B2u
Py·+ 2 B1u
Py·+ 2 Au
Py·+ 2 B1u
PyH·
3
Py 3 A−g
HCO· 2 A”
Pyrene
λc
FWHM
312.7
7.1
325.3
10.0
341.5
14.0
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
Pyrene:H2 O
λc
FWHM
319.2
6.5
329.2
3.2
334.0
4.4
363.2
3.6
354.0
6.5
344.9
6.2
445.6
6.6
435.5
10.2
423.0
12.2
413.8
5.3
490.1
10.0
399.4b
5.2b
...a
...a
...a
...a
405.0b
4.5b
...
...
...
...
...
...
...
...
...
...
...
...
a
Absorption feature was too weak to perform an accurate fit
b
Features are too weak at 25 K; the 125 K values are indicated
Pyrene:CO
λc
FWHM
319.4
7.5
329.2
2.3
334.3
4.1
...a
...a
...a
...a
...a
...a
445.3
7.8
...a
...a
...a
...a
...a
...a
a
...
...a
400.5
4.2
392.5
6.7
378.4
15.7
406.2
4.8
513.4
17.5
535.3
12.5
556.3
14.5
583.0
16.8
604.9
10.0
639.2
15.1
6.3.1 Neutral pyrene bands
As in Chapter 5, the strong, negative bands peaking at 334.0 nm and weaker bands at
329.2 and 319.2 nm in the H2 O ice (see Fig. 6.1), and at slightly shifted positions in the
CO ice, are assigned to the 1 B2u ←1 Ag electronic transition of neutral pyrene (S2 ← S0 )
based on previous studies of pyrene in rare gas matrices [Vala et al. 1994, Halasinski
et al. 2005]. To study the chemistry in absolute number densities, a value of f = 0.33 is
adopted from the literature for the oscillator strength of pyrene [Bito et al. 2000, Wang
et al. 2003]. This value is used throughout this paper both for the Py:H2 O and Py:CO
experiments. Pure pyrene ice measured at 10 K shows broader absorptions located at
341.5, 325.3, and 312.7 nm (see Table 6.1). We did not perform VUV experiments on the
pure pyrene sample.
128
6.3 Band assignments and band strength analysis
6.3.2 Pyrene cation bands
Positive bands at 363.2, 354.0, and 344.9 nm appear upon photolysis in the Py:H2 O experiments. This progression is assigned to the 2 B1u ←2 B3g vibronic transition of the pyrene
cation (Py·+ ) in accordance with the proximity to the band positions reported by Vala et al.
[1994] and Halasinski et al. [2005]. This transition for Py·+ in H2 O ice was reported in
Chapter 5. The 2 B1u ←2 B3g transition is too weak to be detected in the Py:CO experiment.
A stronger Py·+ progression occurs at 445.6, 435.5, 423.0, and 413.8 nm in water ice. Of
these bands, only the strongest at 445.3 nm is detectable in the irradiated Py:CO ice. This
progression is assigned to the 2 Au ←2 B3g transition of Py·+ . The much weaker absorption
caused by the 2 B1u ←2 B3g Py·+ transition at 490.1 nm in H2 O is again undetectable in CO.
In these H2 O and CO ice experiments, Py·+ formation is the result of direct single
photon ionization of the neutral species, following:
VUV
Py −−−→ Py·+ + e− .
(6.2)
Normalized integrated absorbance Py cation
We emphasize that ionization in Py:H2 O ices is far more efficient than in Py:CO ices.
Additional measurements on Py:CO:H2 O mixtures indicate that the presence of H2 O indeed enhances the ionization. Hence, it is possible that water contamination in the CO
ice is responsible for the formation of some, if not all, of the cation species in the Py:CO
0.35
25 K
0.30
50 K
75 K
100 K
0.25
125 K
0.20
0.15
0.10
0.05
0.00
0.0
-0.1
-0.2
-0.3
-0.4
-0.5
-0.6
-0.7
Normalized integrated absorbance Py neutral
Figure 6.2 Integrated absorbance of the 445.6 nm Py·+ band growth plotted against the
loss of the 334.0 nm Py band in the 25, 50, 75, 100, and 125 K ices. Ten seconds elapse
between subsequent data points. Values are normalized to the maximum integrated absorbance of neutral pyrene. The straight line portion of these plots is used to determine
the oscillator strength of Py·+ as described in §6.3.2.
129
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
experiment. The role of water contamination in CO ice will be discussed in more detail
in §6.3.3.
Using baseline corrected spectra as shown in Fig. 6.1, the photochemical evolution
is tracked by integrating areas of bands produced by each species in every spectrum and
plotting them as a function of photolysis time. The strongest Py·+ band at ∼445 nm is
selected to track the number density evolution of this species. To put the kinetic analysis
(§6.4) on a quantitative footing, we determine the oscillator strength of the 445 nm Py·+
band as follows. First, the integrated absorbance of the 445 nm Py·+ band is plotted
versus that of the 334 nm Py band during the course of VUV photolysis at different ice
temperatures. These graphs are shown in Fig. 6.2. It should be noted that there is a
tight, linear behavior between the loss of neutral pyrene and growth of the pyrene cation
during early photolysis times up to 100 s (the first 10 successive datapoints). Inspection
of Fig. 6.2 shows that the slope is steepest and the ratio of the integrated absorbance of the
Py·+ band to the Py band is optimum in the 25 K ice. Since no other photoproduct bands
are evident during the linear correlation stage, we conclude that during this phase, neutral
pyrene is converted solely into the cation as described previously for naphthalene and
quaterrylene [Gudipati & Allamandola 2006a]. The straight-line portion, fitted through
the first 10 data points of irradiation at 25 K, is used to determine the oscillator strength
of Py·+ . Given that the ratio of the Py·+ to the 334 nm Py band is 0.99 and the oscillator
strength of this Py transition is 0.33, the oscillator strength of the 443 nm Py·+ band in
water ice is also taken to be 0.33. This conclusion is consistent with ab initio calculations
on pyrene by Weisman et al. [2005]. They calculated that the oscillator strength of the
cation is only ∼2% stronger than that of the neutral species. As described below, the
photolysis of Py in water ices at higher temperatures produces other species in addition to
the cation. This explains the different curves in Fig. 6.2.
6.3.3 HCO bands in Py:CO
VUV irradiation of a Py:CO ice also produces a vibrational progression ranging from
∼500 to 650 nm. As shown in Fig. 6.3, these absorption bands, located at 513.4, 535.3,
556.3, 583.0, 604.9, and 639.2 nm, are assigned to the 2 A′′ (0, ν′′ , 0) ← X2 A′ (0,0,0) HCO·
(ν”=8–13) transitions based on band positions reported by van IJzendoorn et al. [1983].
The clear HCO· progression indicates a photolytic source of free H atoms in the ice. In
addition, it confirms the ability of this setup to record reactive intermediates in the ice.
A possible explanation of the source of H atoms is related to the experimental conditions. The experiments reported here are performed under high vacuum (10−7 mbar) conditions. Therefore, background H2 O vapor has ample time to condense onto the sample
window while cooling down and growing the ice sample. Water is well known to photodissociate upon VUV irradiation [e.g., Öberg et al. 2009d, Andersson & van Dishoeck
2008] according to
VUV
H2 O −−−→ OH· + H· .
(6.3)
An experiment on VUV irradiation of a “pure” CO ice indicated that HCO· is also effi130
6.3 Band assignments and band strength analysis
ciently produced in the absence of pyrene. Therefore, it is likely that water contamination
is responsible for the production of HCO· via
H· + CO → HCO· .
(6.4)
Another possible formation route could be by means of VUV-induced hydrogen abstraction from pyrene. This pyrene photodissociation reaction, however, is unlikely to occur,
since PAHs are generally highly photostable molecules.
Wavenumber / cm
20000
19000
18000
17000
-1
16000
15000
-3
5.0x10
.
HCO
-3
''= 8
9
2
2
A''(0, '',0)
10
X A'(0,0,0)
11
12
13
Absorbance
4.0x10
-3
3.0x10
-3
2.0x10
-3
1.0x10
0.0
500
520
540
560
580
600
620
640
660
Wavelength / nm
Figure 6.3 Vibrational progression of HCO· generated in a Py:CO ice at 10 K after 600
seconds of VUV irradiation plotted together with a Gaussian fit (red) to the absorption
spectrum (black). The individual Gaussians are shown in blue. The orange line indicates
the red wing of the underlying broad absorption feature (§6.3.6). (This figure is available
in color in electronic form.)
6.3.4 The 400 nm band carrier
Another vibrational progression appears at 400.5, 392.5, and 378.4 nm in the CO ice
experiments. As shown in Fig. 6.4, the 400.5 nm band dominates this progression. In
contrast, a single band appears at 399.4 nm in the Py:H2 O ice upon VUV irradiation of
the samples. The relative intensity of these bands varies with respect to the Py·+ bands.
The 400 nm bands are more pronounced than the cation bands in the H2 O ice only at high
temperatures, whereas they are more pronounced in the low temperature CO ice.
Two additional measurements were performed to identify the carrier responsible for
these transitions. A kinetic experiment was performed on non-VUV-irradiated Py:CO
131
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
6
.
5
Integrated absorbance PyH
.
PyH
4
380
385
390
395
400
405
410
wavelength / nm
3
2
1
Py
0
310
315
320
325
330
335
340
Wavelength / nm
0
-5
-10
-15
-20
Integrated absorbance Py
Figure 6.4 Integrated absorbance of the 400 nm PyH· band growth plotted against the loss
of the 334 nm Py band in a at 10 K CO ice after VUV irradiation is stopped. The straight
line directly reflects the relative oscillator strength of both bands as described in §6.3.2.
The two insets show the PyH· and Py vibrational progressions in a CO ice 90 minutes
after photolysis is stopped.
ice. This ice showed no sign of pyrene ionization by the Xe-lamp, which is used as a
spectroscopic light source.The production of HCO· and the formation of the 400 nm band
were not observed either. Subsequently, the ice was irradiated by the VUV source for
10 minutes. The steady growth of the 400 nm band with VUV photolysis indicates that
the species responsible for the 400 nm band is a product of the VUV processing of the
ice. Moreover, when the VUV irradiation is stopped, the 400 nm band carrier continues
to grow at the expense of the remaining neutral pyrene. This indicates that the chemical
reaction leading to the formation of this species is not directly photon-dependent, but
rather depends on the diffusion of a photoproduct. A similar experiment on a Py:H2 O
ice at 25 K indicates that the same process also takes place in H2 O ice. The detection
of HCO· radicals in the ice and the inherent presence of free photolytic H atoms, implies
that the growth of the vibrational progression starting at ∼400 nm could be the result of
the reaction of pyrene with diffusing H atoms
Py + H· → PyH· .
(6.5)
This assignment to the 1-hydro-1-pyrenyl radical (PyH· ) is supported by other experimental studies [Okada et al. 1976, 1980], where progressions at similar band positions
are observed upon (laser) flash photolysis.
In contrast to the Py:H2 O experiments where pyrene is also efficiently ionized, the
experiment on PyH· formation in CO shows no sign of other reaction products. The
132
6.3 Band assignments and band strength analysis
integrated absorbance of the growing PyH· transition is plotted versus the integrated absorbance of the diminishing neutral in Fig. 6.4. Growth is tracked over a duration of more
than 1.5 hours. Since there is a one-to-one conversion of Py to PyH· in the Py:CO ice
(Eq. 6.5), as described in §6.3.2, we derive an oscillator strength of 0.089 for this species
by fitting a line through the correlating absorbances in Fig. 6.4.
6.3.5 The 405 nm band carrier
Besides the Py·+ and the PyH· bands, another distinct absorption is found in the spectra of
VUV irradiated ices. This feature is located at 405.0 nm in the Py:H2 O and at 406.2 nm in
the Py:CO experiment. In Chapter 5, we tentatively assigned this absorption to a negative
ion, Py− or PyO− . The experiments on Py:CO ices presented here enable us to exclude this
assignment because of the nearly absent Py·+ transitions. Firstly, Py− is ruled out because
a much stronger second Py− absorption band, expected at 490 nm [Montejano et al. 1995],
is absent in our Py:CO experiment. Secondly, PyO− is also ruled out, because it should
exhibit absorption bands down to 350 nm [Milosavljevic & Thomas 2002], bands that are
also absent in the Py:CO experiment. Additionally, in Chapter 5 we assumed that PyO−
was a product of PyOH. The formation of PyO− is also unlikely in the absence of PyOH
absorption in these experiments, as discussed below.
The absorption at 405 nm does not correlate with that of the cation, nor with the
PyH· band. The band only appears during photolysis and hence is characterized as a
VUV-photon-related product. From the literature, it is known that a pyrene triplet-triplet
(3 A−g ←3 B+2u ) transition is expected at this wavelength upon laser excitation of pyrene in
solution which populates the lowest member of the triplet manifold [e.g., Hsiao & Webber
1992, Langelaar et al. 1970]. For the 405 nm band to originate from this triplet-triplet
transition, the lowest level must be populated and remain so with a long enough lifetime
to allow absorption to the 3 A−g level. In the ice experiments reported here, there are a
number of possible routes for pumping the 3 Py state. The most obvious route is by means
of photoexcitation followed by intersystem crossing
1
VUV
isc
Py −−−→ 1 Py∗ −−→ 3 Py.
(6.6)
Triplet formation is found to decrease with decreasing temperature in ethanol ice [Stevens
et al. 1967]. This translates to our experiment in a nearly absent 405 nm band in the low
temperature Py:H2 O experiment, because of the high Py·+ formation efficiency. In the
high temperature H2 O ice experiments, on the other hand, the 405 nm absorption is much
stronger because pyrene is available.
In the CO ice, on the other hand, where Py·+ production is low, formation of the
405 nm band carrier appears to be very efficient at low temperatures. The production of
the 405 nm band carrier requires VUV photons to be initiated. The pumping of the 3 Py
state can again occur by means of Eq. 6.6. Moreover, CO has a dipole allowed electronic
transition in the VUV. Hence, speculating, pumping of the 3 Py state by collisional de133
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
excitation of CO molecules exited by the VUV radiation provides a reaction path of
VUV
CO −−−→ CO∗ ,
(6.7)
CO∗ +1 Py → 3 Py + CO.
(6.8)
followed by
In summary, while we cannot identify the carrier of the 405 nm band, the 3 A−g ←3 B+2u
transition seems a plausible explanation.
6.3.6 Broad absorption feature
Finally, besides the narrower bands reported in the previous paragraphs, we discuss a
broad underlying feature extending from about 350 to 570 nm, which grows upon photolysis in all cases. This band probably comprises overlapping bands caused by a number of
Py/H2 O or Py/CO photoproducts. Part of this Py-residue feature remains even after warming up the sample window to room temperature, whereas all other features disappear at
the water desorption temperature.
As discussed above, the very broad feature must be produced by a variety of similar
but distinct photoproducts, all containing the pyrene chromophore. Mass spectral analysis
of the species produced by the VUV photolysis of a few other PAHs in water ice show
that the parent PAH is not destroyed but that OH, O, and H are added to some of the
edge carbon atoms [Bernstein et al. 1999]. Given the multiplicity of the side sites on
pyrene that can undergo substitution, it is likely that the photoproducts produced in the
experiments reported here are multiply substituted, rather than singly substituted. Thus, it
is plausible that a mixture of related but distinct Py–Xn species, where X may be H, OH,
or O, produce the broad band.
In Chapter 5, we reported the production of a clear and reproducible PyOH band at
344.9 nm in a low temperature H2 O ice. The results presented here do not show evidence
of this absorption feature. However, in some instances the absorption was detected upon
irradiation or warm-up of the ice. The irregular appearance of the PyOH absorption feature in these experiments indicates that the formation of this species is highly sensitive
to the sample’s physical parameters, i.e., structure of the ice, temperature, and concentration. One possible explanation is that in the experiments reported in Chapter 5, the Py
concentration was not controlled and those experiments sampled a very different ice concentration and, by implication, physical ice structure. While we do not have a solution for
this discrepancy, we emphasize that both measurement series have been fully reproducible
over many independent experiments for periods of months. An experimental program to
investigate the role of the PAH:H2 O concentration on ice photochemistry is underway.
134
6.4 Py:H2 O ice photochemistry at different temperatures
6.4 Py:H2 O ice photochemistry at different temperatures
Figure 6.5 shows the spectral evolution of two different Py:H2 O samples at different temperatures. The top frame presents the 280 to 540 nm spectra of the 100 K Py:H2 O ice
after 0, 20, 40, 80, and 160 s of in situ photolysis and the bottom frame the corresponding
spectra for the 25 K ice. These spectra are snapshots of the more than one thousand spectra collected during 4 hr of photolysis. They illustrate the rapid changes that occur during
the early stages in the photochemistry of these ices and the major differences in reaction
products at different temperatures.
To probe the VUV-driven photophysics and reaction dynamics for a set of selected
temperatures, the production and depletion of species was tracked as a function of irradiation time. To this end, the Py 334 nm, Py·+ 445 nm, and PyH· 400 nm bands were
integrated for every spectrum. The spectra in Figs. 6.1 and 6.5 show that it is rather
straightforward to determine the boundaries needed to integrate these bands. We estimate
that the uncertainty in most of these band areas is of the order of 10%.
The integrated absorbances of the neutral Py, strongest Py·+ , and PyH· bands in H2 O
ice at temperatures of 25, 50, 75, 100, and 125 K are plotted versus photolysis time (VUV
fluence) in Fig. 6.6. The spectra in Fig. 6.5 and photochemical behavior in Fig. 6.6 show
that, upon photolysis, neutral pyrene loss is immediate and rapid. The initial growth of
Py·+ mirrors the rapid, initial loss of Py. However, while Py steadily decreases, and several
other Py photoproduct bands increase during some 4 hours of photolysis, the production
of Py·+ reaches a maximum and then slowly diminishes. From Fig. 6.6, one can clearly
see that ionization of pyrene is most efficient in the low temperature ice. Formation of
PyH· , on the other hand, is far more efficient at higher temperatures.
For comparison, the integrated absorbances for the irradiated Py:CO ice are plotted
as a function of time in the right bottom frame of Fig. 6.6. It should be noted that the
PyH· band is multiplied by a factor of 10 in the Py:CO experiment, compared to a factor
of 20 in the Py:H2 O experiment. The PyH· band is clearly more prominent in the CO
ice experiment than in the H2 O ice experiments. The Py+ signal on the other hand is
negligible. This indicates that the H2 O ice plays a role in ion formation and stabilization.
To place this behavior on a quantitative footing, the integrated areas for the Py and
Py·+ bands are converted to number densities using Eq. 6.1. Here, an oscillator strength
of 0.33 is used for the 334 nm Py bands. The values used for the oscillator strengths of
the Py·+ and PyH· bands are 0.33 and 0.089, respectively, as determined in §6.3.2 and
§6.3.4. Perusal of Fig. 6.6 shows that Py behaves similarly in all of the ices considered
here. Regardless of temperature, its signal drops quickly with the onset of irradiation
and continues to diminish with ongoing photolysis. Likewise, Py·+ grows rapidly with
initial photolysis but peaks after a relatively short time interval corresponding to a fluence
of roughly 8 × 1016 photons and then drops continuously. While the Py·+ growth and
loss curves resemble each other, cation production efficiency is strongest in the 25 K ice.
This efficiency remains of the same order at even lower temperatures (not shown here).
The photolysis time required for the cation to reach a maximum shortens with increasing
135
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
3
Wavenumber /10 cm
34
32
30
28
26
24
-1
22
20
0.02
0.00
T = 100 K
-0.02
0 s
-0.04
20 s
40 s
Optical depth (-)
-0.06
80 s
160 s
-0.08
-0.10
0.10
0.08
0.06
0.04
0.02
0.00
-0.02
T = 25 K
-0.04
-0.06
-0.08
-0.10
300
350
400
450
500
Wavelength / nm
Figure 6.5 The VUV-induced spectroscopic changes in Py:H2 O ice for two different temperatures as a function of photolysis time. Comparing the spectra from the 25 K ice
(bottom) with those of the 100 K ice (top) shows the critical role that temperature plays
in determining photochemical pathways in a PAH-containing ice. In the 25 K ice, cation
formation is favored over production of the pyrene residue and the 400 and 405 nm band
carriers. The opposite holds for the 100 K ice.
136
6.4 Py:H2 O ice photochemistry at different temperatures
Photon fluence / 10
0
200
400
600
800
1000
1200
1400
0
15
photons
200
400
600
800
1000
1200
1400
1.0
T = 25 K
T = 50 K
0.8
Error bar
Py
Py+
0.6
PyH (x20)
Fit Py
Fit Py+
0.4
Fit PyH
Products
Normalized integrated absorbance (-)
0.2
0.0
1.0
T = 75 K
T = 100 K
T = 125 K
T=10 K (Py:CO)
0.8
0.6
0.4
0.2
0.0
1.0
Py
0.8
Py+
PyH (x10)
0.6
0.4
0.2
0.0
0
2000
4000
6000
8000
10000
12000
14000
0
2000
4000
6000
8000
10000
12000
14000
Time / s
Figure 6.6 The integrated absorbance of the Py 334 nm, Py·+ 445 nm, and PyH· 400 nm
bands as a function of VUV irradiation in Py:H2 O ices at 25, 50, 75, 100, and 125 K
and a Py:CO ice at 10 K plotted together with the fits (grey lines) described in §6.4.
Integrated absorbance values are scaled and normalized to the initial value for the Py
signal. The PyH· feature is multiplied by a factor of 20 for the Py:H2 O experiments and
by a factor of 10 for the Py:CO experiment. The approximate, overall growth of the total
Py photoproduct band (P1 +P2 +P3 ) is also shown (dotted line). (This figure is available in
color in the on-line electronic form.)
137
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
temperature. The PyH· band contribution is minor with respect to the Py·+ band for ices
below 50 K. This reverses between 50 and 75 K, suggesting that there is a change in the
dominant Py:H2 O ice photochemical channel in this temperature range.
PyH.
k3
P3
k21
k22
Py
k1
P1
k11
k12
.
Py +
k2
P2
Figure 6.7 Reaction scheme used to fit the experimental data.
A kinetic analysis of the plots in Fig. 6.6 is carried out using the reaction scheme
indicated in Fig. 6.7. Here, k11 is the photoionization rate of Py to Py·+ , k12 the electronion recombination rate of Py·+ , k21 the production rate of the PyH· feature, and k22 the rate
of the reverse reaction of PyH· to Py. The rates designated k1 , k2 , and k3 are the production
rates for the different products that comprise the Py-residue band. The oscillator strengths
for the Py·+ and PyH· bands are also fitted, but are restricted to remain within ±10% of
the experimentally determined values of 0.33 and 0.089. All reactions are assumed to be
first order in the reactant. The relative abundances of “free or solvated electrons” and O,
H, and OH radicals in the ice are not considered.
The fits to the growth and decay curves are included in Fig. 6.6 and the temperature
dependence of the derived rate constants is presented in Fig. 6.8. The agreement between
the fit and the experimental data in terms of curve shape and absolute intensity is good.
The fitted oscillator strengths of the Py·+ and PyH· bands amount to 0.31 and 0.082,
respectively, and hence do not deviate much from the experimentally determined values.
The graph in Fig. 6.8 indicates that the Py photoionization rate (k11 ) drops rapidly between 25 and 50 K. The electron recombination rate (k12 ) decreases only slightly, if at all,
within the errors over the entire temperature range. As mentioned above, the production
of the PyH· becomes more important at higher temperatures. Its formation rate (k21 ) is low
in all ices up to 50 K (< 4.4 × 10−5 ), but jumps to > 1×10−4 in the ices with temperatures
of 75 K and higher. The back channel from PyH· to Py, k22 , also shows a temperature
dependence. It increases almost linearly in going from cold to warm ices. The formation
rate of a photoproduct produced directly from Py (k1 ) also seems to jump at 50 K. The
formation rate of products originating in the Py·+ species, on the other hand, seems to
lower with increasing temperature. Finally, the rate of product formation from the PyH·
channel is low throughout the entire temperature range. The jump in rate of the formation
of P1 and PyH· with temperature probably reflects the diffusion barrier of radical species
(H· and OH· ) in the ice.
Since published studies of the processes induced by the photolysis of other PAH:H2 O
ices are limited, not much information is available with which to compare these results.
While, to the best of our knowledge, there are no reports of the photochemistry that takes
138
6.4 Py:H2 O ice photochemistry at different temperatures
-3
-3
2×10
2×10
k11
k12
-3
-3
1×10
-1
Rate coefficient (s )
0
1×10
0
100
50
0
150
-4
0
100
50
150
-3
2×10
1×10
k21
k22
-4
-4
1×10
0
5×10
0
100
50
-3
150
0
0
1×10
4×10
k1
-4
0
k2
-4
50
100
150
0
k3
-6
3×10
0
150
-6
6×10
5×10
100
50
-4
2×10
0
50
100
150
0
0
50
100
150
Temperature (K)
Figure 6.8 Parameters (knm ) as a function of temperature resulting from fitting the reaction
scheme (Fig. 6.7) to the kinetic experiments (Fig. 6.6). All rates are indicated in s−1 .
place as a function of ice temperature or of long-term fluence, the VUV photochemistry
of the PAHs naphthalene, 4-methylpyrene (4MP), and quatterrylene in water ice at 10 K
has been studied [Gudipati & Allamandola 2003, Gudipati 2004, Gudipati & Allamandola 2006a,b]. The results obtained are in good agreement with the low temperature
(25 K) case reported here. Namely, the parent PAH is easily and efficiently ionized, by
quantitative conversion of the neutral species to the cation form. The focus of the earlier
studies was on cation production and stabilization and not on long-duration photolysis
experiments. In their study of 4MP:H2 O (1:>500) ice at 15 K, Gudipati & Allamandola
[2003] utilized a reaction scheme similar to that on the right half of that presented in
Fig. 6.7. Table 6.2 compares the reaction rates that they determined with those of the
25 K ice reported here. Except for the production of P2 , which differs by one order of
magnitude, there is very good agreement between the rate constants for each step in the
two experiments.
The growth and decay curves in Fig. 6.6, taken together with the temperature depen139
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
Table 6.2 The reaction rates for the VUV photolysis of Py:H2 O (∼1:5,000) ice at 25 K
compared to those for 4-methylypyrene:H2 O (1:≤500) ice at 15 K [Gudipati & Allamandola 2003].
Rate
VUV
Py−−−→Py·+
Py·+ +e− →Py
Py → P1
Py·+ → P2
a
(k11 )
(k12 )
(k1 )
(k2 )
This work
(s−1 )
Photon ratea
(cm2 photon−1 )
Gudipati 2003
(s−1 )
(1.2±0.1)×10−3
(9±2)×10−4
(5±1)×10−5
(5±1)×10−4
1.2×10−17
9×10−18
...
...
1.3×10−3
8×10−4
4×10−5
5×10−5
Photon rates are indicated only for reaction channels which are dominated by photon processes.
dence of the reaction rates in Fig. 6.8, show that the VUV-driven PAH photochemistry
depends strongly on ice temperature. The influence of the ice morphology on this chemistry was also investigated, to understand the origin of the temperature dependence. An
experiment on an ice deposited at 25 K, annealed to 125 K, and subsequently cooled to
25 K before photolysis, showed that the ionization rate and efficiency are similar to that
of an unannealed ice. Apparently, it is not the morphology but the temperature of the
ice that primarily determines which process dominates. We discriminate between two
temperature regimes. One governed by ion-mediated processes that dominate at 25 K
and slightly higher temperatures, and a second, presumably radical-driven regime, that
becomes increasingly more important at higher temperatures.
6.5 Astrochemical Implications
As shown in the previous paragraphs, ionization and chemistry of a rather small PAH
species, pyrene, trapped in H2 O ice turns out to be very efficient in a laboratory setting.
Here, we extend these findings to interstellar conditions, with the aim of including the
calculated rates in astrochemical models. For this, it is crucial to distinguish pure photochemical processes from diffusion, since the latter will be highly dependent on the number
density of radicals and electrons in the ice. As mentioned in the previous paragraph, the
photoionization of Py is probably a single-photon process, whereas protonation of Py and
the electron recombination of Py·+ are the results of both VUV photolysis and diffusion.
The mechanism for PyH· deprotonation is unclear, since it can proceed by means of either VUV processing or through hydrogen abstraction by diffusing species. Diffusion of
radicals through the ice is a thermally activated process and will therefore increase with
temperature. Recombination, however, is largely temperature-independent in our experiments, indicating that the rate of Py·+ recombination is not dominated by the diffusion
of electrons in the ice. If Py·+ loss occurs by means of electron recombination and not
Py·+ reaction with H2 O or one of its photoproducts, the electron most likely originates
from the initial photoionization event after which electrons remain in the vicinity of the
recombining Py·+ species. Hence, this local process can be, although indirectly, regarded
as a single-photon process.
140
6.5 Astrochemical Implications
The rates of protonation of Py and deprotonation of PyH· show a temperature dependence and the importance of diffusion can therefore not be excluded. This makes it harder
to directly translate the rates (s−1 ) into photon rates (cm2 photon−1 ). However, we can
determine astrochemical photon rates for both ionization and recombination of pyrene in
interstellar H2 O ice (see Table 6.2).
Now, to translate this to the astrochemical situation and with other processes, we
assume that PAHs generally have an ionization rate similar to that of pyrene. How do
ionization and chemistry compare with other processes such as the photodesorption of
the icy grain mantle, in which the PAHs are embedded? To exemplify this, the rate of
ionization of a PAH in water ice at 25 K (in photon−1 ) is calculated anywhere in a dense
cloud where AV = 3 and compared with the VUV photodesorption rate of H2 O derived
by Öberg et al. [2009d]. It is well established that the onset of ice formation occurs in
clouds with an edge-to-edge (through the cloud) magnitude of AV = 3 [e.g. Whittet et al.
2001]. Thus, inside our hypothetical dense cloud at AV = 3 (from cloud edge to within
the cloud), ices are present.
The experimentally determined PAH ionization rate in H2 O at 25 K, normalized to the
total amount of deposited PAH is given by
+
k11 =
]
d [PAH
[PAH]0
dt
= 10−3 s−1 .
(6.9)
Consider a typical interstellar grain, covered by a 100 monolayer (ML) thick ice. The
number of sites on a grain is 1015 cm−2 . If we assume that one in every 104 particles
on the grain is a PAH, the total number of PAH molecules on the grain is [PAH]0 =
100 × 1015 × 10−4 = 1013 cm−2 . Furthermore, the VUV photon flux in our laboratory,
Φ, is 1014 photons cm−2 s−1 . The production rate of PAH cations on an interstellar grain
is now given by [PAH]0 · k11 /Φ = 10−4 photon−1 . This ionization rate is an order of
magnitude lower than the rate of photodesorption (∼ 10−3 photon−1 ) [Öberg et al. 2009d].
However, in our dense cloud the number of photons available for PAH photoionization is larger than the number of photons available for photodesorption of H2 O ice. This
is because H2 O photodesorption is primarily caused by VUV photons, whereas PAH ionization can occur for much lower energy photons. To quantify the radiation field in a
dense cloud at AV = 3 as a function of wavelength (λ), we take the average UV interstellar radiation field (Iν ) from Sternberg [1988] and rewrite the expression to Iλ with units
photons cm−2 s−1 nm−1
Iλ =
1.068 × 10−4 c 1.719 × 10−2 c 6.853 × 10−1 c
−
+
,
λ3
λ4
λ5
(6.10)
where c is the speed of light in nm s−1 . The attenuation of the radiation field by dust as a
function of wavelength is given by
#
−11.6AV Aλ
,
Dλ = exp
RV
A1000Å
"
(6.11)
141
6 Pyrene:H2 O ice photochemistry: ion-mediated astrochemistry
from Draine & Bertoldi [1996], where we assume that RV = 3.1 and AV /A1000Å = 0.21
[Whittet 2003]. This results in
"
#
Aλ
Dλ = exp −0.8 AV ,
(6.12)
AV
where the table of Aλ /AV values is taken from Mathis [1990]. The photon flux per second
per wavelength interval is given by
Pλ = Iλ Dλ .
(6.13)
Water ice absorbs photons with wavelengths ranging from 130 to 150 nm [Kobayashi
1983, Andersson & van Dishoeck 2008]. The ionization energy of PAHs on the other
hand, is lowered by about 2 eV when in H2 O ice [Gudipati & Allamandola 2004, Woon
& Park 2004]. For the wavelength range available for ionization of PAHs, assuming that
H2 O blocks all photons below 150 nm, we take 150 to 250 nm [Li & Draine 2001]. By
integrating the photon flux in a cloud of AV = 3 over both wavelength intervals a number
of photons available for PAH ionization is found that is 6 times larger than the number of
photons available for photodesorption of H2 O. Additionally, at AV = 3, the cosmic-rayinduced UV field is negligible compared to the interstellar UV field [Shen et al. 2004].
Therefore, the occurrence of photoionization is of similar order as photodesorption of the
main component in the grain mantle in a dense cloud. The ionization rates from Table 6.2
can be directly included in astrochemical models in the form
d[PAH+ ]
= k11 Ψ[PAH],
dt
(6.14)
where [PAH+ ] is the concentration of the PAH (pyrene) cation in the ice, k11 is the photon rate in cm2 photon−1 , Ψ is the photon flux in photon s−1 cm2 , and [PAH] is the PAH
(pyrene) concentration in the ice.
In the above calculation, we assume that all PAHs exhibit the ionization behavior of
the pyrene chromophore. Of course, more PAHs need to be investigated experimentally
before drawing conclusions on their general photochemical behavior in interstellar ices.
However, if all PAHs have an ionization rate similar to that of pyrene, photoionization
and subsequent chemical reactions of PAHs trapped in ices are important processes in
dense clouds. When frozen-out in ices, PAHs have an important impact on the radical and
electron budget in solid state chemistry. Hence, the processes described here may be more
important than previously assumed in modeling complex interstellar grain chemistry.
6.6 Conclusions
A recently constructed setup has been used to track, on a sub-second timescale, the photochemistry of a PAH in H2 O and CO ices as a function of temperature. The setup used
here clearly has advantages compared to relatively slow infrared photochemical ice studies. The conclusions from this work on a PAH, pyrene, trapped in an interstellar ice
analogue are summarized below:
142
6.6 Conclusions
1. A set of photochemical reaction products has been identified in both irradiated
Py:H2 O and Py:CO ice experiments. The reaction products result from direct photoionization of pyrene, or from a reaction of the parent, pyrene, with free H atoms
produced in the matrix. Additionally, an absorption band is tentatively assigned to
a triplet-triplet transition of pyrene. A vibrational progression assigned to HCO· is
found in spectra of the VUV-irradiated Py:CO ice.
2. Pyrene is easily and efficiently ionized when trapped in H2 O ice. Photoionization is
a non-diffusion-related reaction and hence a photon rate of 1.2×10−17 cm2 photon−1 ,
which can serve as input for astrochemical models, is derived.
3. When trapped in CO ice, pyrene ionization is inefficient compared to that in water
ice.
4. Electron-ion recombination is independent of ice temperature and is characterized as a non-diffusion-dominated reaction. For this process, a photon rate of
9 × 10−18 cm2 photon−1 is derived, which can be directly used in astrochemical
models.
5. There are two distinct reaction paths in the photochemistry of pyrene trapped in
H2 O ice. At low temperatures (< 50 K), the chemistry is dominated by ionmolecule interactions and processes. At temperatures above 50 K, reactions are
dominated by diffusing radical species.
6. A simple model indicates that, in dense clouds where AV = 3, the rate of pyrene
ionization is comparable to the rate of photodesorption in water-rich ices. Hence,
chemical reactions involving pyrene and its cation can be important in modeling
grain chemistry in these environments.
143
CHAPTER 7
Ionization of Polycyclic Aromatic Hydrocarbons
trapped in H2O ice 1
Mid infrared emission features originating from Polycyclic Aromatic Hydrocarbons are
present throughout many phases of the interstellar medium. Towards dense clouds, however, these features are heavily quenched. Observations of dense clouds point out that
many simple molecules are frozen out on interstellar grains, forming thin layers of ice.
It is likely that more complex non-volatile species, such as PAHs, also freeze out on
grains and contribute to the chemistry of interstellar ices. The study presented here aims
at obtaining reaction rate data for the photochemistry of PAHs in an interstellar H2 O ice
analogue. Furthermore, the experimental data are implemented in a chemical model of a
dense interstellar cloud in order to study the relevance of PAH:H2 O ice reactions in these
interstellar regions. Time dependent near-UV/visible spectroscopy on anthracene, pyrene,
benzo[ghi]perylene and coronene containing interstellar H2 O ice is performed at 25 and
125 K, using an optical absorption setup for the study of ices (OASIS). Near-UV/VIS absorption spectra are obtained for these four PAHs and their cationic species trapped in
H2 O ice. Relative oscillator strengths of the cation absorption bands are derived relative to the oscillator strength of the neutral parent PAH. The number density evolution of
species in the H2 O matrix is measured and fitted to a reaction scheme, resulting in rate
constants for the corresponding reactions. A freeze-out model is employed to determine
on what timescale PAH molecules are incorporated in interstellar ices. The PAH:H2 O
photochemical rate constants are used in an astrochemical model, which is used to determine the importance of PAH:H2 O ice photoprocessing in going from a dense cloud to
a protostellar object. All four PAHs studied here are found to be readily ionized upon
VUV photolysis when trapped in H2 O ice and exhibit similar rates for ionization. The
PAH freeze out occurs on rather long time scales in a dense cloud. Thus, PAH photoprocessing will only be important after the PAH containing ices are formed, i.e. during the
protostellar phase. In this phase, photoprocessing of PAH containing H2 O ice is indeed
an effective process.
1 Based on: J. Bouwman, H. M. Cuppen, M. Steglich, L. J. Allamandola, H. Linnartz, Astronomy and
Astrophysics, in prep.
145
7 Ionization of PAHs in interstellar ices
7.1 Introduction
The presence of Polycyclic Aromatic Hydrocarbons (PAHs) in many phases of the interstellar medium is evidenced by their strong and ubiquitous mid-infrared emission features
[Smith et al. 2007, Draine et al. 2007, Draine & Li 2007, Tielens 2008]. Mid-IR features
are efficiently emitted by a PAH after being excited by an energetic photon. Toward
dense clouds, however, the mid-IR features are strongly quenched. Here, most volatile
molecules are frozen out on grains forming layers of ice [e.g., Pontoppidan et al. 2004,
Boogert et al. 2008, Öberg et al. 2008, Bottinelli et al. 2010]. Under such conditions,
PAHs most likely also condense on interstellar grains, incorporating them in interstellar
ices.
Experimental studies on the effect of vacuum ultraviolet (VUV) irradiation of interstellar ice analogues have shown that more complex molecules can be formed in the simplest mixed ices [e.g. Gerakines et al. 1995, Öberg et al. 2009c]. The first laboratory
studies on VUV irradiated PAH containing ices indicated that PAHs are easily ionized.
These experiments also show the formation of new species. Time dependent information
on these chemical reactions, however, is largely lacking.
Here, we present the time evolution of the destruction of four PAHs, anthracene (Ant,
C14 H10 ), pyrene (Py, C16 H10 ), benzo[ghi]perylene (Bghi P, C22 H12 ), and coronene (C24 H12 )
in H2 O ice together with the formation and destruction of the ionized PAH+ species. This
chapter aims to quantify and understand the time dependent chemistry of PAH:H2 O ice
mixtures upon VUV irradiation and the resulting photoproducts.
The chemical evolution is tracked by means of near-UV/VIS absorption spectroscopy
at two extreme temperatures, 25 K and 125 K. This work is an extension of the detailed
study of pyrene in H2 O ice presented in Chapter 4 and aims to draw more general conclusions on the PAH photochemistry in ices based on a larger sample set. Furthermore, the
present study extends the PAH:H2 O photochemistry to larger and astrophysically more
relevant members of the PAH family.
The outline of this paper is as follows. In §7.2 the experimental setup is briefly discussed, together with the details of the theoretical calculations. Paragraph 7.3 describes
spectra of the PAH and PAH+ cations and present their (relative) oscillator strengths and
the assignments of the observed transitions. The fitted time dependent data are discussed
in detail in §7.4, after which the astrophysical implications are discussed in §7.5. Finally,
the conclusions are summarized in §7.6.
7.2 Experimental technique
Here we briefly describe the experimental setup. The system is described in detail in
Chapter 5. The setup consists of a high-vacuum (∼10−7 mbar) chamber. In the center of
the vacuum chamber a MgF2 sample window is suspended, which is cooled by a closed
cycle He cryostat to a temperature of 25 K. Temperatures as low as 11 K can be realized.
PAH containing H2 O ices are grown onto the sample window by vapor deposition. MilliQ H2 O is further purified by three freeze-pump-thaw cycles and the PAHs are used as
146
7.2 Experimental technique
commercially available (Ant, Aldrich ≥99%, Py, Aldrich, 99%, Bghi P, Aldrich, 98%, Cor,
Aldrich, 99%). The thickness of the samples is monitored by laser interference and the
amount of deposited PAH is monitored in absorption, allowing for determination of the
samples PAH:H2 O concentration.
The inlet system has been modified for measuring the large PAHs in our sample, Bghi P
and Cor. A sample container is mounted in the vacuum chamber and located adjacent
to the H2 O deposition tube. The sample container is heated with polyimide insulated
nichrome heater wire. The H2 O flow to the sample is set such that a certain static pressure
is reached inside the vacuum chamber, and the current through the heater wire is chosen
such that the rate of deposition results in the desired sample concentration. Additionally,
a heat shield is mounted on the sample holder, such that the PAH that vaporizes during
heating of the sample container to the desired temperature is collected on the heat shield,
rather than on the sample window.
After deposition on the 25 K window, the sample is heated to the desired temperature.
Subsequently, the sample is subject to Vacuum Ultra-Violet (VUV) radiation, which is
produced by a H2 flow microwave (MW) discharge lamp. The lamp operates at a static
H2 pressure of 0.4 mbar, and a MW power of 100 W, resulting in an effective VUV flux
of ∼ 1014 photons·cm−2 s−1 at the sample surface.
Near UV/VIS absorption spectra are taken during VUV processing of the samples.
To this end, a Xe-lamp is used as a broadband light source and a spectrometer equipped
with a 1024×256 pixel CCD camera is used as detector. The CCD camera is read out
in vertical binning mode by a computer on which the raw data are converted into optical
depth (OD = ln(I/I0 )). Spectra ranging from ∼280 to 800 nm are taken at a rate of 0.1 s−1 ,
which is sufficient to monitor chemical changes in our ice samples. Each spectrum is the
result of co-adding 229 individual spectra, resulting in an excellent signal-to-noise
ratio.
R
The integrated absorbances of the deposited neutral PAH signal, τν dν, is converted
into a PAH column density, NPAH , via the oscillator strength of the neutral PAH, f , by:
R
τν dν
.
(7.1)
NPAH =
8.88 × 10−13 f
Together with an ice thickness measurement based on the interference pattern in the reflection of a HeNe laser as described in Chapter 5, this allows for a rather accurate determination of the PAH:H2 O concentration. Sample deposition is stopped at thicknesses of
∼2 µm, resulting in comparable ice samples.
In a typical experiment of 4 hours, as many as 1400 spectra are obtained. The spectra
are all baseline corrected by fitting a second order polynomial through points where no
absorptions occur and subsequently subtracting the polynomial. Additionally, absorption
band are integrated and, if necessary, corrected for contributions by atomic H-lines originating in the H2 MW discharge lamp. All the data handling is performed in a LabView
program.
In order to support the assignments of measured absorption bands caused by the
photo-products, we performed density functional theory (DFT) calculations using the
gaussian09c software [Frisch et al. 2009]. We used the B3LYP functional in conjunction
147
7 Ionization of PAHs in interstellar ices
with the 6-311++G(2d,p) basis set to determine the ground state geometry and electronic
structure of PAH neutrals and cations. Excited states were investigated within the framework of the time-dependent density functional theory (TDDFT) applying the same level
of theory.
7.3 PAH:H2 O spectroscopy
Long duration photolysis experiments are performed on a set of four PAHs (Ant, Py, Bghi P,
and Cor) in H2 O ice at low (25 K) and high (125 K) sample temperatures. For determining
the amount of deposited PAH, oscillator strengths for the neutral PAHs are taken from the
literature. The amount of deposited H2 O is measured by laser interference, yielding ice
thicknesses of typically ∼2 µm. Combining the thickness of the sample and the amount
of deposited PAH results in the PAH concentration in the sample. An overview of the
used mixture concentrations , the temperature at which the samples are photolyzed, and
the oscillator strengths values ( f ) of the neutral PAH adopted from the literature is given
in Table 7.1.
Table 7.1 An overview of the studied PAH species and used PAH:H2 O concentration,
sample temperature as well as the wavelength interval of the strongest neutral absorption
band system, and the corresponding literature values for the oscillator strengths.
Species
Ant
Ant
Ant
Ant
Py
Py
Bghi P
Bghi P
Cor
Cor
Conc.
1:2.000
1:1.500
1:900
1:500
1:5.000
1:6.500
1:2.000
1:1.000
1:5.000
1:4.000
TSample (K)
25
25
25
125
25
125
25
125
25
125
λrange (nm)
316–381
f
0.1a
295–350
0.33b
320–388
0.21c
273–314
1.04d
a
Gudipati [1993] b Bito et al. [2000], Wang et al. [2003] c Rouillé et al. [2007] d Ehrenfreund et al. [1992]
For each of the PAHs under investigation, we determined the oscillator strengths of
the cationic species’ absorption bands relative to that of the neutral precursor. This allows
for a full quantitative study of the formation and destruction, which will be presented in
§7.4. Relative oscillator strengths are determined by plotting the time evolution of the
integrated absorbance of the cation transition under investigation against the integrated
absorbance of the strongest electronic transition of the neutral species (see Fig. 7.1). Both
integrated absorbances are normalized to the amount of deposited neutral PAH, which is
determined during the preparation of the ice sample. A quantitative conversion of the PAH
molecule into its cationic species is assumed to occur during the first photolysis stage:
148
7.3 PAH:H2 O spectroscopy
VUV
PAH −−−→ PAH+ + e− .
(7.2)
Linear fits are made to the first data points, of which the slope directly reflects the
oscillator strength of the cation relative to that of the neutral. None of the PAH species
in our sample substantially deviates from a one-to-one conversion during the first 100 s
of photolysis, making the assumption valid and the resulting relative oscillator strength
values reliable vantage points for further analysis.
Relative integrated absorbance cation (-)
0.6
Ant
0.5
Py
BghiP
Cor
0.4
0.3
0.2
0.1
0.0
0.0
0.1
0.2
0.3
0.4
0.5
0.6
Relative integrated absorbance neutral (-)
Figure 7.1 The correlation between the amount of produced cation and the amount of used
up neutral, both relative to the total deposited amount of neutral PAH.
Typical absorption spectra for the 25 K PAH:H2 O ice mixtures taken at the maximum
cation absorption are plotted in Fig. 7.2. The position of the band origin, the range used
for integration and oscillator strength value relative to the strongest neutral absorption are
listed in Table 7.2. The assignments of the neutral and photoproduct bands is discussed
for each individual PAH below.
7.3.1 Anthracene (C14 H10 )
The negative signal between ∼310 and 380 nm is caused by the destruction of the neutral
Ant molecules and reflects the depopulation of the 1 B2u ← 1 Ag transition of neutral Ant
[Bak et al. 2000]. The positive absorption features throughout the spectrum are caused
by species that are produced by photodestruction of the parent PAH. A strong vibronic
progression arises between 500 and 760 nm with its maximum at 719.6 nm. This progression has previously been assigned to the 2 Au ← 2 B2g transition of the singly ionized
Ant species (Ant+ ) [Szczepanski et al. 1993b]. For this transition we derive an oscillator
149
7 Ionization of PAHs in interstellar ices
Frequency / 103 cm-1
35
30
+
Ant
25
+
Ant
+
Py
Optical Depth
20
17.5
15
12.5
+
Ant
+
Py
a
+
Py
b
+
BghiP
+
BghiP
+
BghiP
c
Cor
+
Cor
+
Cor
+
d
300
350
400
450
500
550
600
650
700
750
800
Wavelength / nm
Figure 7.2 The 280 to 800 nm spectra of the PAHs anthracene (a), pyrene (b),
benzo[ghi]perylene (c), and coronene (d) in H2 O ice, photolyzed at 25 K. Negative features indicate that a species is destroyed, positive bands indicate that a species is formed.
The mixing ratios are 1:700, 1:5.000, 1:2.500, and 1:4.000 (PAH:H2 O) for anthracene,
pyrene, benzo[ghi]perylene, and coronene, respectively. The molecular structures are
also indicated in the corresponding spectra.
150
7.3 PAH:H2 O spectroscopy
strength value of 0.59 relative to that of the neutral. A sharp and strong absorption feature previously assigned to the 2 B1u ← 2 B2g transition of Ant+ appears at 351.1 nm. This
absorption has an oscillator strength of 0.15. Likewise, an absorption feature which is
assigned to the 2 Au ← 2 B2g transition of Ant+ is found at 313.7 nm.
Table 7.2 Overview of the studied PAHs, state symmetry, position of the band origin, the
range over which the transition is integrated, and oscillator strength of the cation species
relative to that of the strongest neutral transition.
Species
Ant
Ant+
Ant+
Ant+
Py
Py+
Py+
Py+
PyH·
Bghi P
Bghi P+
Bghi P+
Bghi P+
Cor
Cor
Cor+
Cor+
Cor+
a
Symmetry
1
B2u (a)
2
Au (b)
2
B1u (b)
2
Au (b)
1
B2u (c,d)
2
B1u (c,d)
2
Au (c,d)
2
B1u (c,d)
...
1
B2 (e)
2
B1 ( f )
2
B1 ( f )
...
1
B1u
1
E1u
B1,2g ( f )
B1,2g ( f )
(f)
Origin Pos. (nm)
375.4
719.6
351.1
313.7
334.0
363.2
445.6
490.1
399.4
379.8
762.2
509.7
404.3
337.6
301.4
687.1
463.7
362.5
range (nm)
316-385
505-753
349-354
307-318
290-345
350-370
411-470
...
380-410
320-389
720-788
451-533
390-410
320-341
276-311
630-760
389-473
352-370
frel.
1.00
0.74
0.15
0.37
1.00
0.13
0.99
...
0.26
1.00
0.13
1.10
0.13
0.17
1.00
0.20
0.23
0.16
b
Bak et al. [2000] Szczepanski et al. [1993b] c Halasinski et al. [2005] d Vala et al.
[1994] e Rouillé et al. [2007] f indicates a tentative assignment based on theoretical calculations presented here.
Two more absorptions are apparent in the spectra of our photolyzed sample. One
sharp absorption appears around 445.8 nm and is probably due to photolysis of small Py
contaminations in our ice sample, resulting in a Py+ absorption. Additionally, a broad
feature spanning the range from 380 to 470 nm is found. This band does not correlate
with the cation features and is hence thought to be caused by a mixture of Ant+H and/or
Ant+OH addition reaction products. These reaction products have previously been mass
spectroscopically identified in VUV photolyzed Ant:H2 O (1:≥100) mixtures [Ashbourn
et al. 2007].
7.3.2 Pyrene (C16 H10 )
The VUV photolysis of Py:H2 O mixtures has been studied and is described in detail by
Chapter 6. Here we only shortly describe the band assignments.
151
7 Ionization of PAHs in interstellar ices
The negative bands that appear between 290 and 345 nm are assigned to the 1 B2u ←1 Ag
electronic transition of neutral pyrene (S2 ← S0 ) [Vala et al. 1994, Halasinski et al. 2005].
Most of the positive bands that form upon VUV photolysis of the Py containing H2 O ice
are ascribed to the Py+ species. The system ranging from ∼411-470 nm is the strongest
Py+ transition and is assigned to the 2 Au ←2 B3g transition. The weaker absorption bands
between 350 and 370 nm are assigned to the 2 B1u ←2 B3g Py+ vibronic transition. Finally,
the band on the red-wing of the strongest Py+ transition is due to the 2 B1u ←2 B3g transition. Besides the rather strong Py cation absorptions, two more bands are detected around
400 and 405 nm. The band at 400 nm was previously found to originate from an electronic transition in PyH· and the band at 405 nm was tentatively assigned to an electronic
transition of 3 Py.
7.3.3 Benzo[ghi]perylene (C22 H12 )
The negative bands between ∼280 and 390 nm in the spectrum of the irradiated Bghi P:H2 O
indicate that the neutral species is destroyed upon VUV photolysis. The absorption bands
have previously been assigned to the S2 (1 B2 ) ← S0 (1 A1 ) transition of Bghi P [Rouillé et al.
2007]. In turn, new bands appear upon VUV photolysis of the Bghi P containing H2 O ice.
A very strong absorption, which has previously been assigned to a Bghi P+ absorption by
Salama et al. [1995] arises at 509.7 nm. Another, much weaker absorption appears in
the mid-IR at 762.2 nm. This band has been assigned in previous matrix work to a low
energy electronic transition of the Bghi P+ species [Hudgins & Allamandola 1995a]. Here
we report an absorption band at 404.3 nm, which shows a clear correlation with the other
Bghi P+ absorptions. We ascribe this band to a higher energy electronic transition of the
Bghi P+ species.
An attempt has been made to assign the new observed cation transitions. The optimized geometry of the Bghi P cation is found to be of C2v symmetry. The calculations are
based on the molecule in the x-z plane with the z axis coinciding with the C2 symmetry
axis. The electronic ground state is 2 A2 making dipole-allowed transitions to A2 , B2 , and
B1 states possible. In the observed wavelength range, several transitions are predicted
by the TDDFT calculations. A transition to a 2 B1 state is calculated to be at 673 nm
(fcalc. =0.048), which could give rise to the observed band at 762.2 nm. The next strong
transition is found at 472 nm (fcalc. =0.21) relatively close in energy to the strongest observed band at 509.7 nm and we tentatively attribute the corresponding transition to 2 B1 .
The calculations also predict a transition to the 2 A2 state around 300 nm, which overlaps strongly with an absorption from the neutral molecule and consequently cannot be
assigned experimentally.
7.3.4 Coronene (C24 H12 )
The neutral Cor molecule is of D6h symmetry. From the A1g ground state, dipole-allowed
transitions are only possible to electronic states of A2u or E1u symmetry. The observed
152
7.4 PAH ionization rates
absorption spectrum in cryogenic matrices strongly resembles the spectrum of hexa-perihexabenzocoronene [HBC, Rouillé et al. 2009]. The weak S1 (1 B2u ) ← S0 (1 A1g ) transition
is not seen in our spectrum, but the S2 (1 B1u ) ← S0 (1 A1g ) transition appears at 337.6 nm.
Like in HBC, it gains intensity due to vibronic interaction with the first allowed transition
S3 (1 E1u ) ←S0 (1 A1g ) found at 301.4 nm. The TDDFT calculations predict this transition at
303 nm (fcalc. =0.65). The first two transitions are predicted to appear at 386 and 363 nm.
As already noted by Oomens et al. [2001], the degenerate ground state of the coronene
cation causes Jahn-Teller interaction and leads to an effective reduction of the point group
from D6h , as found for the neutral species, to D2h . It also complicates the assignment of
measured absorption bands on the basis of DFT calculations. However, the DFT geometry optimization predicts a slightly elongated structure for the coronene cation with an
elongation of only 0.7% of the total diameter. Therefore, we can only provide tentative
band assignments assuming D2h to be the correct point group. In that case, the electronic
ground state is 2 B3u . The strongest feature seen in H2 O ice at 463.7 nm could correspond
with a transition to an excited state of B1,2g symmetry. This band was found at 462 nm
in an Ar [Szczepanski & Vala 1993] and at 459 nm in a Ne matrix with an f-value of
0.012 [Ehrenfreund et al. 1992]. The calculation predicts three states between 390 nm
and 460 nm with somewhat higher oscillator strengths between 0.02 and 0.05. Likewise,
the broad cation absorption at 687.1 nm could belong to states of B1,2g symmetry. Further
dipole-allowed transitions to B1,2g electronic states are predicted around 310 nm, overlapping with the strongest absorption of neutral Cor. These bands are probably very broad,
leading to a raise in the baseline around 310 nm as visible in Fig. 7.2.
7.4 PAH ionization rates
In the previous section we derived the oscillator strengths of the cation bands relative to
those of the neutral parent PAH. These numbers are now used for quantification of the
reaction channels which are involved in the VUV photolysis of PAHs in interstellar ices.
In the analysis presented here, the evolution of the column density in the ice is tracked as
a function of time. Cation relative oscillator strengths are used to convert the integrated
absorbance in a column density relative to the amount of neutral. The analysis is based on
the strongest cation bands listed in Table 7.2. The resulting time evolution of the number
densities relative to the deposited amount of neutral PAH is shown in Fig. 7.3 for the four
systems studied here.
In the analysis we consider a channel for ionizing the PAH with rate k11 , a backchannel for recombination of PAH+ species with electrons with rate k12 , a channel for
the formation of products P1 directly from the parent neutral PAH with rate k1 , and the
formation of products P2 from the PAH+ species with rate k2 . The reaction scheme is
schematically displayed in Fig. 7.4. For the Py:H2 O sample, a reaction scheme with one
more channel was used in Chapter 6 since this molecule clearly follows an additional
reaction path involving PyH· , which can be unambiguously tracked spectroscopically.
The contribution of PyH· to the total amount of reaction products is low and for the sake
of analysis and comparison all data, including pyrene, are fitted with the reaction scheme
153
7 Ionization of PAHs in interstellar ices
15
Photon fluence / 10
0
200
400
600
800
1000
1200
1400
0
photons
200
400
600
800
1000
1200
1400
8000
10000
12000
14000
1.0
PAH
25 K
125 K
PAH+
fit to PAH
0.8
fit to PAH+
0.6
0.4
0.2
0.0
1.0
0.8
Relative column density (-)
0.6
0.4
0.2
0.0
1.0
PAH
PAH+
PAH fit t dep.
0.8
PAH+ fit t dep.
PAH fit t indep.
PAH+ fit t indep.
0.6
0.4
0.2
0.0
1.0
0.8
0.6
0.4
0.2
0.0
0
2000
4000
6000
8000
10000
12000
14000
0
2000
4000
6000
Photolysis time / seconds
Figure 7.3 The PAH neutral decays and rise and fall of the corresponding cation signal
for four PAHs, anthracene, pyrene, benzo[ghi]perylene, and coronene for two different
temperatures (25 K and 125 K). The molecule structures and fitted curves are indicated in
the plots.
154
7.4 PAH ionization rates
indicated in Fig. 7.4 only, i.e. omitting the formation and destruction of PyH· .
PAH
gas
kacc
k11
PAH
H
OH
P
e
grain
k1
1 grain
-
PAH
k12
k2
P
.
+
grain
H
OH
2 grain
Figure 7.4 The reaction scheme used to fit the experimental time evolution of PAHs and
their cations upon VUV irradiation is indicated in the dotted box. The total reaction
scheme including the accretion of PAH species from the gas phase into the ice is used for
modeling the astrophysical case in §7.5.
The time dependent chemistry of PAHs in H2 O ice is studied for two temperatures,
25 K and 125 K. Fits to the data are co-plotted with the experimental data in Fig. 7.3.
Fits to the time evolution curve of the PAH with the strongest cation absorption, Bghi P, are
made twofold; keeping the ionization channel temperature dependent and independent.
This yields insight in the effect of the temperature on the process. In Chapter 6 we noted
a temperature dependence in the Py ionization channel. The fits to the data, however,
are very sensitive to the integrated cation signal. In the case of Py+ , the signal is very
weak and thus an accurate fit is hard to obtain. From the two fits to the Bghi P data in
Fig. 7.3 it is clear that temperature actually does not have a large influence on the quality
of the fit, and inherently the ionization rate k11 turns out to be independent of temperature.
The ionization rate k11 of the other species, Ant and Cor, also does not exhibit a large
temperature dependence. Table. 7.3 gives an overview of the fit parameters which are
obtained while keeping all parameters free, i.e. dependent of temperature.
From the fit data constants in Table 7.3 it is clear that the recombination channel, k12 ,
increases with temperature, except for the case of Cor. Additionally, the rate of product
formation directly from the parent PAH, k1 , seems to be independent of temperature. The
rate of formation of products from the cation species, k2 , on the other hand, drops to zero
for all PAHs except for Bghi P. What is most striking about the data presented in Table 7.3,
is that the ionization rate for all PAHs is of the same order.
At the end of an experiment, both the neutral parent PAH and the cation features
are destroyed. There are absorptions superposed on the baseline, but there is no clear
spectral signature which can be compared to literature data on possible photoproducts.
Paper I in this series employed mid-IR spectroscopy to track the spectral changes in VUV
155
7 Ionization of PAHs in interstellar ices
Table 7.3 Fit parameters to the experimental time evolution of the destruction of neutral
PAHs and formation of cation species in PAH:H2 O ice following the reaction scheme
indicated in Fig. 7.4. All reaction rates are in units of 10−4 s−1
Species
Ant
Ant
Py
. Py
Bghi P
Bghi P
Cor
Cor
Temp (K)
25
125
25
125
25
125
25
125
k11
8
7
10
8
7
4
10
9
k12
6
20
8
50
2
20
20
20
k1
4
6
4
8
0.8
2
5
7
k2
0.8
0
3
0
1
1
2
0
photolyzed ices of somewhat higher concentration. In that paper, more clear signatures of
photoproducts were found, which were tentatively assigned to fundamental vibrations of
functional groups in the newly formed photoproduct species (PAH–Xn , with X being H,
O, or OH). It is likely that similar photoproducts are formed in the experiments described
here.
7.5 Astrophysical implication
As indicated by the rate constants in Table 7.3 in the previous section, PAH chemistry
and PAH ionization are rather efficient processes in VUV irradiated PAH containing H2 O
ice. This was already known for Py, but is shown here also to be the case several other
PAHs. The largest molecule in our sample, Cor, contains 24 carbon atoms, which is
still small from an astrophysical viewpoint. However, the experimental study presented
here indicates that the rate of ionization is rather size independent in the range of species
investigated here. We therefore extend our findings to the astrophysical case, in which
larger PAH species (NC ≥ 50) are though to be most relevant.
Here, we discuss the relevance of PAH:H2 O ice chemistry in translucent or dark cloud
conditions (A ≥ 2). Whittet et al. [2001] did not observe H2 O ice in clouds with an
edge-to-edge visual extinction of A ≤ 3. We assume that this roughly corresponds to an
edge-to-center visual extinction of A ≤ 1.5 and hence at least H2 O ice will be able to exist
under the translucent or dark conditions considered here.
In the interstellar medium, PAHs are initially in the gas phase. The formation of PAH
photoproducts on grain mantles therefore consists of two steps: first the neutral PAHs
freeze out from the gas phase onto the grains where they can then participate in the solid
state reaction network as schemetically indicated in Fig. 7.4. The rate of accretion of PAH
156
7.5 Astrophysical implication
species onto the grain, RISM
acc , is given by
2
RISM
acc = vPAH ngrain πa nPAH
r
r
8k ngrain 2 T gas
=
πa
nH nPAH
π nH
M
r
T gas
nH nPAH
= 4.57 · 10−8
M
= kacc nH nPAH ,
(7.3)
with vPAH the velocity of the PAH molecule in the gas phase, nPAH the gas phase number
n
density of PAHs, nH the total number density of hydrogen, ngrain
the dust to gas number
H
−12
ratio (10 ), a is the standard grain radius (0.1 µm), T gas the gas temperature, M the
molecular mass of the PAH molecule (amu), and kacc the accretion rate (cm3 s−1 ). Additionally, all measured PAH reaction coefficients, k, are scaled to the interstellar photon
flux and the extinction of the cloud by
!
ΨCR
ISM
lab ΨISM
k
=k
exp −γAV +
,
(7.4)
Ψlab
Ψlab
with klab the measured ionization rate constant (s−1 ), ΨISM the interstellar UV flux (photons cm−2 s−1 ), Ψlab the laboratory UV flux (ΨISM =10−7 Ψlab ), γ a measure of UV extinction relative to visual extinction (≈2) [Roberge et al. 1991], AV the visual magnitude, and
ΨCR the cosmic ray induced photon flux.
An initial total gas phase PAH abundance of 4% with respect to H2 O and an abundance
of H2 O of 10−4 with respect nH is assumed. We further use the largest PAH investigated
in this study, Cor, as a prototype system, which results in M = 300 amu, leaving T gas , nH ,
and AV as input parameters for the model.
The timescales at which PAHs freeze out onto the grains are investigated first. The
left panel of Fig. 7.5 plots the gas phase PAH abundance for different initial densities nH .
The graph clearly shows that for a dense cloud with nH = 103 cm−3 , it takes more than 107
years for the gas phase to become depleted of PAHs. Densities of 105 cm−3 and higher
need to be reached will PAHs freeze out on a reasonable timescale. The reason for this is
that at low densities the frequency with which PAHs encounter a grain is very low, since
the grain abundance directly scales with density. Furthermore, the accretion rate scales
with the velocity of the species and inherently with its mass. PAH molecules are heavy
molecules and thus move slow through the cloud, thereby slowing down their depletion
process. Once the PAH does encounter a cold grain, the sticking probability is high and
since PAHs are highly non-volatile molecules, they will remain in the ice as long as the
ice matrix remains to exist.
The problem is that at such high densities, the interstellar radiation field is almost
fully attenuated and only cosmic ray induced photons play a role. This is not enough
to get a high processing rate of the PAHs within a reasonable time. In Chapter 4 we
have shown for the high mass Young Stellar Object (YSO) W33A and the low mass YSO
157
7 Ionization of PAHs in interstellar ices
0
10
2
10
7
nH = 10
PAH
-1
Relative PAH surface abundance
-3
Gas phase PAH abundance (cm )
10
1
10
6
nH = 10
0
10
5
nH = 10
-1
10
4
nH = 10
-2
10
3
nH = 10
-2
PAH
+
10
-3
10
P1
P2
-4
AV = 2
10
0
10
PAH
-1
10
-2
10
PAH
+
P1
P2
-3
10
-3
10
-4
10
2
10
4
10
Time (yr)
6
10
AV = 3
2
10
4
10
6
10
Time (yr)
Figure 7.5 Left panel: Modeled depletion of PAHs on cold interstellar grains as a function
of cloud density. Right top panel: Modeled photoprocessing of condensed PAH species in
a cloud with a visual magnitude of AV =2. Right bottom panel: Modeled photoprocessing
of condensed PAH species in a cloud with a visual magnitude of AV =3.
RNO 91 that up to 3% of the ice mantle may consist of PAH photoproducts. These high
PAH photoproduct concentrations must be formed under the influence of a rather strong
radiation field. The two right plots in Fig. 7.5 show the PAH ice chemistry as a function
of time under influence of the standard interstellar radiation field at AV = 2 (right top) and
3 (right bottom). The model starts will PAHs frozen out onto the grains. The neutral PAH
and the photoproduct abundances are given with respect to the total PAH abundance. For
AV = 2 the onset of the formation of photoproducts is after 104 yrs; for AV = 3 the onset
occurs after 105 yrs.
In summary, to explain the high abundances of frozen out PAH photoproducts reported
in Chapter 4, grains first need to be in a high density environment after which they are
exposed to a high UV field. This corresponds to the following scenario for both the high
and low mass YSO’s W33A and RNO 91: the PAHs:H2 O ices form in the pre-collapse
high density phase. Once the newly formed star starts radiating UV photons, some of the
grains will be exposed to the UV field. In a protoplanetary disks a similar situation can
158
7.6 Conclusions
occur due to vertical mixing. The ice covered dust grains form in dense regions around
the midplane. Vertical mixing brings them closer to the warmer top layer of the disk and
back down towards the cold mid-plane of the disk.
7.6 Conclusions
The work presented here describes photolysis experiments on interstellar H2 O ice samples
containing four different PAHs. The experiments are performed for two temperatures;
one low (25 K) and one high (125 K) temperature. Experimental data are fitted to a
reaction model and the resulting rate constants are used in a first order astrochemical
model, including freeze-out of species and photoprocessing of ices. The conclusions are
summarized below:
1. Near UV/VIS spectroscopy on the photolysis of four PAH species (Ant, Py, Bghi P,
and Cor) trapped in an interstellar H2 O ice is performed. Photoproduct bands have
been assigned to electronic transitions of PAH cation transitions of the respective
PAH species.
2. The temporal evolution of the production of cation bands is tracked for the four
PAHs at the two temperatures under investigation. Oscillator strengths of the PAH+
species have been derived for all the PAH+ electronic transitions relative to those of
the neutral parent PAH molecule. Derived relative oscillator strengths of the PAH+
transitions are used to quantify the temporal evolution of species.
3. It is found that all four PAHs behave similarly upon VUV photolysis in a H2 O
ice. The cationic species is efficiently produced in the temperature ices, the number
density reaches a maximum and then slowly subsides. The ioniziation efficiency
is decreased upon photolysis of PAHs in high temperature H2 O ice. As concluded
in Chapter 4 and 6, this behavior can be attributed to PAH-radical or PAH+ -radical
reactions being more important due to a larger mobility of radical species (H+OH)
in the ice at these temperatures.
4. The experimental PAH and PAH+ column density time evolution data have been
fitted with a model based on a chemical reaction scheme involving PAH ionization
and PAH reactions with radical species. Rate constants are derived and reported.
All four PAHs exhibit similar reaction rates, allowing for the general conclusion
that PAH photoreaction rates are rather size-independent over the range of species
studied here.
5. A model to calculate the freeze-out of PAHs on cold interstellar grains in a dense
molecular cloud indicates that PAH depletion is rather ineffective on short time
scales, because of the PAHs mass and number density. On time scales of the formation of a protoplanetary disk, however, PAHs have efficiently frozen out in ices.
159
7 Ionization of PAHs in interstellar ices
6. The photochemistry of PAH:H2 O ices is modeled for an a protostar. The model
results point out that the onset of PAH:H2 O photoprocesing occurs at t = 104 yrs
for a visual magnitude of AV = 2 and at 105 yrs for AV = 3. Thus, photoprocessing
of PAHs in ices is expected to be of importance in more evolved objects.
160
CHAPTER 8
Future challenges
Astrophysical laboratory techniques have been and will be of unprecedented value for
interpreting and guiding astronomical observations. Spectroscopic data, both in the gas
phase and in the solid state, allow to identify species in space and to derive inter- and
circumstellar abundances. Experimentally derived rate constants, in addition, serve as
input for astrochemical models which can be directly compared to observations. In the
last decades, the progress in laboratory based research has boosted the understanding of
chemical processes in space. With the further improvement of new upcoming observational facilities and the refinement of chemical models, there is a need for more and more
detailed laboratory data. In this chapter experimental challenges are addressed that will
be useful for astronomy using OASIS, the new setup that has been described here.
In Chapters 5 – 7 it was shown that it is possible to study the photochemistry of PAHs
in water ice upon VUV irradiation through optical (i.e. electronic) spectroscopy using
direct absorption spectroscopy. The method allows studying reaction products in situ and
in real time. As such the technique offers a very versatile and generally applicable tool
that is capable of studying other systems as well, both substantially larger and smaller.
This provides information needed for identifying species in space as well as insight in
possible reaction schemes. The latter point has been addressed in detailed in this thesis,
but the use of this work for an astronomical detection of a PAH in the solid state, has only
been mentioned in §5.5. The optical PAH ice data, presented in this thesis, basically hold
the promise to search for optical solid state PAH signatures in space, as an alternative to
electronic gas phase work that has been unsuccessful so far. A kick-off project with the
ESO-VLT equipped with the FORS2 spectrometer towards the embedded K0 star DoAr21
has been performed to search for broad absorption features that may be correlated to the
absorption bands described for the PAHs and PAH-derivatives in this thesis. The analysis
of the data is still in progress and at this stage only the effort to identify PAH in space
through their ice spectrum is reported.
It is generally expected that also (and according to several publications, particularly)
large PAHs containing 50 carbon atoms or more, the so called GRAND-PAHs, and their
photoproducts are present in space. These are expected to be formed in the stellar ejecta
of dying stars. Gas phase spectra of such complex PAHs are lacking, mainly because of
the experimental challenges that go along with bringing such molecules in the gas phase
161
8 Future challenges
in the laboratory. The only data available on large PAHs today are from matrix isolation
spectroscopy. In such experiments, mainly argon and neon are used as a molecular surrounding, as for such environments matrix interactions are as small as possible. OASIS
can be used in a similar way, as a regular matrix setup, guiding gas phase studies. Test
experiments have been performed on Bghi P (C22 H12 ) and hexa-peri-hexabenzocoronene
(HBC C42 H18 ) using argon instead of water as a matrix. The results are shown in Figure 8.1. The distortion of the PAH energetic structure is proportional to the polarizability
of the matrix material. If the purpose of the experiments would be to measure the energy
of free PAH molecules, i.e. gas phase species, the experiments are to be performed for
more than one matrix material. One could for example measure the spectra in xenon,
argon and neon. If the transition energies are now plotted as a function of polarizability
of the matrix, it is possible to extrapolate the energy level to zero polarizability, i.e. the
electronic energy in vacuum. This is a good indication of the origin of gas phase electronic transitions. This method, subsequently, can be employed to select (GRAND-)PAHs
which possibly contribute to diffuse interstellar band absorption features.
35
30
Frequency /10
25
3
20
cm
-1
15
Absorbance (-)
40
250
300
350
400
450
500
550
600
650
700
750
800
Wavelength / nm
Figure 8.1 The spectrum of matrix isolated Hexa-peri-benzo-coronene and its cation in
argon at 12 K (bottom) plotted together with matrix isolated benzo[ghi]perylene and its
cation band in argon at 12 K (top)
The general applicability of OASIS is also demonstrated with a set of test measurements on a different type of carbon containing molecules, the nano-diamonds adamantane
(C10 H16 ) and diamantane (C14 H20 ). The species are abundantly present in meteorites and
are thought to exist in other regions of the interstellar medium as well. Calculations predict that these species are easily ionized by Lyman-alpha radiation and that their cationic
species have moderately strong electronic transitions in the optical part of the electromag162
netic spectrum [M. Steglich, priv. comm.]. Again, a possibility exists that such species
contribute to absorptions in the diffuse interstellar medium. The test measurements performed so far, were done in high concentration matrices and did not exhibit absorption
bands that originate from diamondoid ions. Further work is needed.
The next step is to include ‘real’ prebiotical molecules in the ice. In the last year it was
shown that VUV irradiation of a pure methanol ice results in the generation of more complex species such as ethanol, methylformate, acetic acid, glycol aldehyde and ethylene
glycol [Öberg et al. 2009c]. That study was based on a long research tradition in Leiden
and follows the experiments by Greenberg and coworkers who irradiated astronomical ice
equivalents with a VUV broadband light source for many days and identified gas chromatographically amino acids in the resulting organic refractory [Muñoz Caro et al. 2002].
Even though the experiments were performed under quite rough experimental conditions
(the residue had to be analyzed outside the vacuum and after warm-up) this has set the
tone for this field. An alternative way for this bottom-up approach is a top-down scenario
in which the focus is not on the generation of more complex species, starting from simple
precursors, but to include real biological systems — e.g. nucleobases, ribonucleotides
and their biosynthesis precursors — in the ice and to study their photostability. This can
be ideally studied with the new experimental setup and will be a research topic in the next
years.
Another interesting subject that can be addressed with OASIS is that of the production
of small radicals in interstellar ices. In Chapter 6, for example, the observation of HCO·
is reported. OH is another interesting object, as the production rate of OH radicals in
an interstellar water ice is not yet fully understood. As infrared techniques are generally
slow, such methods are not suited to address this question, and the optical equivalent
described here offers an alternative. The first test measurements on irradiated pure H2 O
ices show that some new very broad absorptions indeed present. This is overcome by
doping an argon matrix with H2 O, with a resulting concentration of about 1:100 (H2 O:Ar).
A vibronic progression ascribed to the A2 Σ+ (ν=0)←X 2 Π3/2 (ν=0) transition at 308 nm is
found [Hancock & Kasyutich 2004]. These test measurements can be extended to more
astronomically relevant ices. Ultimately, the diffusion of radical species in the matrix can
be studies as a function of temperature.
Finally, OASIS was initially designed as CESSS: a Cavity Enhanced Solid State Spectrometer. The setup is equipped with two stable and micrometer adjustable mirror mounts
on the in and outlet port of the vacuum system. The initial idea was to use an incoherent broad band cavity enhanced detection scheme [Fiedler et al. 2003] to perform the
experiments that have been described in this thesis. This did not work in the solid state,
presumably because of the light refraction in the ice. Fortunately, a singly pass experiment turned out to be sensitive enough. In the gas phase, however, the system works
fine.
A set of mirrors with reflectivities as high as 99.95% can be aligned perpendicular
by a HeNe laser. The resulting system behaves like a cavity, which can be excited by
intense broad band white light originating from the Xe-lamp. Light will continuously leak
out of exit side of the cavity. This light has traversed through the cavity for some time,
depending on the reflectivity of the mirrors, resulting in an enhancement of the absorption
163
8 Future challenges
Frequency / cm
14500
14450
14400
14350
Absorbance (-)
14550
-1
687
688
689
690
691
692
693
694
695
696
697
Wavelength / nm
Figure 8.2 The absorption spectrum of the b1 Σ+g (ν=1)←X 3 Σ−g (ν=0) oxygen transition
path length, which directly reflects the increase in sensitivity of the detection technique
[Fiedler et al. 2003]. The power of this arrangement is exemplified by a spectroscopic
measurement on the doubly forbidden The forbidden b1 Σ+g (ν=1)←X 3 Σ−g (ν=0) gas phase
oxygen transition can only be observed in the 50 cm long absorption cell, if significant
enhancement is achieved in the cavity [O’Keefe & Deacon 1988]. A typical absorption
spectrum taken under atmospheric pressure is shown in Fig. 8.2. The example indicates
the power of the setup for performing spectroscopic studies on gas phase species. Of
course, this technique can be used for measurements on astrophysically relevant species.
With minor adjustment of the system, one can sensitively perform spectroscopy on a
(plasma) expansion over a large frequency range with moderately high resolution.
In conclusion, OASIS has turned out to be a very versatile tool that has proven its use
for astronomical research.
164
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172
Nederlandse samenvatting
De spectroscopie en de chemie van interstellaire ijs
analogen
Het onderzoek dat in dit proefschrift wordt beschreven richt zich op laboratorium experimenten die astrochemische processen nabootsen onder omstandigheden, zoals die in de
ruimte voorkomen. Deze omstandigheden zijn extreem. In donkere interstellaire wolken kan de temperatuur zakken tot slechts 10 graden boven het absolute nulpunt, zo’n
-263◦ C. Ook de ‘lucht’druk is uitermate gering. Met slechts enkele moleculen per kubieke centimeter zijn interstellaire wolken bijzonder leeg. Lange tijd werd aangenomen,
dat de dichtheden in de ruimte te laag zijn voor chemische reacties. De kans dat een
aantal deeltjes gelijktijdig botst is immer zeer gering. In 1937 werd echter het eerste molecuul in de ruimte ontdekt: CH, een stabiele binding van een koolstof (C) en waterstof
(H) atoom. Deze ontdekking luidde het begin in van een nieuwe onderzoeksdiscipline:
de astrochemie. Vandaag, bijna driekwart eeuw later, staat de teller van in de ruimte geïdentificeerde moleculen boven de 150. Hieronder bevinden zich alledaagse substanties
— water (H2 O), koolstofdioxide (CO2 ) en alcohol (CH3 CH2 OH) —- maar ook meer exotische stoffen, zoals geladen moleculen (HCO+ en C6 H− ) en het grootste tot nu toe in de
ruimte gedetecteerde molecuul: HC11 N (de recentelijke publicatie van C60 in de ruimte is
daarbij nog even buiten beschouwing gelaten). De wijze waarop deze moleculen ontstaan
is maar deels begrepen.
De chemische evolutie van moleculen in de ruimte volgt in grote lijnen het stervormingsproces. Een interstellaire wolk (ruwweg 99% gas en 1% stof) stort onder zijn eigen
gewicht ineen. De kern van de wolk is ‘donker’, en bij gebrek aan licht koelt de wolk sterk
af. In koude gebieden in de ruimte waar zich zowel gas als stofdeeltjes bevinden, kunnen
gas moleculen vast vriezen op het oppervlak, waardoor laagjes ijs ontstaan. In de afgelopen twee decennia is gebleken dat dit ijs, als reservoir voor moleculen en als katalysator,
een belangrijke rol speelt bij de vorming van nieuwe moleculen. Chemische reacties worden daarbij geïnitieerd door bestraling met energetisch licht (harde ultraviolette (UV) of
173
Nederlandse samenvatting
kosmische straling) en door een beschieting met atomen (H, N, C, O, S) of electronen en
daarbij vormen nieuwe (en meer complexe) moleculen. In een volgend stadium ontstaan
in gebieden met nog hogere dichtheid proto-stellaire kernen waaruit uiteindelijk sterren
worden geboren. Rondom de nieuwe ster hoopt zich restmaterie op, gas, stof en ijs, in een
zogenaamde proto-planetaire schijf. Deze wordt vervolgens met UV licht van de jonge
ster bestraald en thermisch verhit, waarbij andere reacties in het ijs plaats vinden. Tevens
verdampen ijs moleculen en komen weer in de gas fase terecht waar ze verder kunnen
reageren. Het stof en gas van de proto-planetaire schijf vormt uiteindelijk het materiaal
waaruit kometen en planeten ontstaan. De chemische evolutie van moleculen in de ruimte
is derhalve ook bepalend voor de chemische samenstelling van planeten.
In de ruimte komen moleculen alleen voor in de gas fase of in de vaste stof en zoals
hierboven beschreven wisselwerken beide fases, door vastvriezen (accretie) en verdampen (desorptie). De nadruk in dit proefschrift ligt op de studie van inter- en circumstellair
ijs. In het laboratorium worden daartoe onder hoog vacuüm condities laagjes ijs gegroeid
die bijvoorbeeld met speciale lampen worden bestraald om het harde ultraviolette stralingsveld in de ruimte na te bootsen.
Figuur 1 Een typisch spectrum van infrarood licht genomen richting een ster in wording.
Uit de ontbrekende kleuren kan afgeleid worden welke moleculen zich in het ijs bevinden.
Vanaf de Aarde (of vooral met telescopen die rond de Aarde draaien) kunnen we deze
dunne ijslaagjes onderzoeken door naar het infrarode licht te kijken, dat door ijs in de
ruimte wordt geabsorbeerd. Bepaalde kleuren licht blijken te ontbreken en het resulte174
Nederlandse samenvatting
rende spectrum is typisch voor het voorkomen van bepaalde chemische substanties in het
ijs, die bij die specifieke kleuren licht (lees energie) gaan trillen. Door deze astronomische spectra te vergelijken met laboratorium data zoals beschreven in dit proefschrift, is
het vervolgens mogelijk om de samenstelling van deze materie op grote afstand te onderzoeken. In Figuur 1 is een infrarood spectrum van een ster in wording weergegeven,
waaruit de samenstelling van het ijs kan worden afgeleid. Spectroscopische details (precieze frequenties, band breedtes en intensiteitsverhoudingen) bieden vervolgens verder
inzicht in de structuur en temperatuur van het ijs, of en hoe het ijs is gemengd, en welke
chemische reacties in het ijs plaats hebben gehad. Dit wordt in dit proefschrift besproken
voor ijs bestaande uit H2 O:CO in hoofdstuk 2 en voor NH3 en CH3 OH houdend ijs in
hoofdstuk 3.
In de afgelopen jaren is duidelijk geworden, dat het waarschijnlijk is dat complexe
(organische) moleculen in ijs worden gevormd. Dit gebeurt ook bij zeer lage temperaturen. Recentelijk onderzoek, bv. van met ultraviolet licht bestraald puur methanol ijs laat
zien, dat vele van de grotere moleculen die al zijn geïdentificeerd in de ruimte, op deze
wijze kunnen worden gevormd. Het is echter nog niet duidelijk hoe complex de chemie in
interstellair ijs daadwerkelijk kan worden, of bv. de bouwstenen van het leven - aminozuren - in de vaste stof kunnen ontstaan. Dit zal in de komende jaren, o.a. in Leiden, worden
onderzocht. Wat wel duidelijk is, dat grote complexe moleculen kunnen vastvriezen in ijs,
en dit is een ander onderwerp dat in dit proefschrift wordt besproken.
1
3
2
4
Figuur 2 Polycyclische Aromatische Koolwaterstoffen waaraan in dit proefschrift metingen zijn verricht: 1) anthracene, 2) pyrene, 3) benzo[ghi]perylene en 4) coronene.
Naast het absorberen van infrarode straling, kunnen bepaalde moleculen ook, nadat ze
door een energetisch lichtdeeltje zijn aangeslagen, infrarood licht uitzenden. De kleuren
van deze emissie zijn net als de kleuren waarbij een molecuul licht kan absorberen, uniek
en dus molecuul specifiek. In 1973 werden sterke infrarode emissiebanden in de ruimte
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ontdekt bij kleuren die specifiek zijn voor koolstof en waterstof houdende moleculen, de
zogenaamde PAKs: Polycyclische Aromatische Koolwaterstoffen. Deze stabiele moleculen worden op Aarde vooral gevormd bij verbrandingsprocessen en zijn bekend om hun
negatieve invloed op de luchtkwaliteit. Het is waarschijnlijk dat PAKs ook in de ruimte
worden gevormd, bijvoorbeeld als bijproduct in verbrandingsreacties in een ster. Dit is in
overeenstemming met de gevonden emissie spectra. Typische voorbeelden van PAKs zijn
in Figuur 2 weergegeven.
Wanneer PAKs in de ruimte voorkomen, dan is het waarschijnlijk dat zij, net als kleinere moleculen, vastvriezen op stofdeeltjes en bijdragen aan de chemische processen die
in het ijs plaats kunnen vinden. Het is slechts deels mogelijk om deze processen met infrarode straling te onderzoeken (zie hoofdstuk 4). De reden hiervoor is, dat PAKs veelal
vergelijkbare vibrationele bewegingen uitvoeren en daarom spectroscopisch moeilijk te
onderscheiden zijn. Een mogelijke oplossing biedt echter ultraviolette en optische spectroscopie, waarbij de electronische eigenschappen van een molecuul worden bestudeerd
en deze zijn voor verschillende PAKs diverser. Derhalve is in de afgelopen jaren een
nieuwe opstelling gebouwd — OASIS (Optical Absorption Setup for Ice Spectroscopy
— (zie hoofdstuk 5), waarmee het mogelijk is PAKs in water ijs met zichtbaar licht te
bestuderen (hoofdstukken 6 en 7). Metingen zijn bovendien niet alleen meer molecuul
specifiek, maar ook sneller en gevoeliger, waardoor het mogelijk is chemische processen
ter plekke vrijwel in ‘real-time’ te volgen.
Dit proefschrift
Dit proefschrift beschrijft een aantal laboratorium experimenten aan inter- en circumstellaire ijs analogen. De metingen zijn verricht met een drietal opstellingen. In Leiden met
een Fourier Transform infrarood spectrometer (zie hoofdstukken 2 en 3) en een optische
spectrometer (zie hoofdstukken 5, 6 en 7) die beide ijs in een hoog vacuüm opstelling
doorlichten en bij NASA Ames met een vergelijkbaar infrarood systeem (hoofdstuk 4).
Bij alle drie opstellingen is het mogelijk het ijs met hard UV licht te bestralen. Het infrarode onderzoek is vooral gericht op kleinere moleculen en het optische werk laat met
name werk aan PAKs zien. Het doel van het onderzoek is om de spectroscopische vaste
stof signatuur te bepalen voor verschillende ijs composities en condities, ook na UV bestraling, zodat astronomische waarnemingen kunnen worden geïnterpreteerd. Verder is
gekeken naar mogelijke reactie schema’s in het ijs en daarbij is vooral optische spectroscopie ingezet.
Hoofdstuk 1 geeft een algemeen overzicht en is een inleiding voor het onderzoek dat
in dit proefschrift is beschreven. In hoofdstuk 2 wordt de wisselwerking tussen twee
belangrijke ijs componenten in de ruimte — CO en H2 O — in detail beschreven. Het
onderzoek beschrijft de veranderingen in de absorptie profielen van zowel CO als H2 O
bij het inmengen van de partner en voor verschillende temperaturen. Dit geeft inzicht
in de interacties in het ijs en biedt verder een mogelijkheid om infrarode astronomische
data te interpreteren. In hoofdstuk 3 is dit expliciet uitgewerkt voor metingen aan methanol (CH3 OH) en ammoniak (NH3 ) in water ijs en spectra zijn vergeleken met recentelijk
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Nederlandse samenvatting
verkregen data van de Spitzer infrarood ruimte telescoop. In dit hoofdstuk wordt de hoeveelheid ammoniak ijs bepaald die zich bevindt in het ijs rond jonge sterren van lage
massa. Verder duiden spectroscopische details aan dat het methanol zeer waarschijnlijk
niet met water is gemengd, hetgeen aangeeft dat beide stoffen een verschillende chemische evolutie hebben doorlopen. Het methanol bevindt zich wel in een CO rijke ijslaag
en dit stemt overeen met het idee, dat CH3 OH ontstaat uit waterstof atoom addities in
CO ijs. Hoofdstuk 4 breidt deze metingen uit naar PAKs in water ijs en laat zien dat na
UV bestraling vele nieuwe moleculen worden gevormd. Een groot aantal nieuwe banden
is gevonden en deels toegekend, vooral aan geïoniseerde PAKs. De metingen zijn vervolgens gebruikt om astronomische spectra te verklaren die gemeten zijn in de richting
van jonge sterren. De belangrijkste conclusie is dat PAKs een deel van de karakteristieke
absorptie in het 6.2 micrometer gebied kunnen verklaren.
Het tweede deel van dit proefschrift beschrijft optische ijs spectroscopie. De nieuwe
opstelling, OASIS, wordt gedetailleerd in hoofdstuk 5 beschreven. De nadruk ligt op de
experimentele details en deze worden geïllustreerd aan de hand van metingen aan pyreen
(C16 H10 ) in zeer koud water ijs. Wanneer het pyreen met hard ultraviolet licht wordt bestraald, verliest het gemakkelijk een electron; het pyreen wordt geioniseerd. De nieuwe
opstelling maakt het mogelijk om dit proces tijdsopgelost te bestuderen. Hoofdstuk 6 gaat
vervolgens in op de chemische reacties die plaats vinden in een pyreen houdend water ijs,
wanneer het met UV licht wordt beschenen. Het ionisatie gedrag en andere chemische
reacties zijn gevolgd voor verschillende temperaturen en de resulterende temperatuursafhankelijkheid wordt besproken. Het blijkt dat de snelheid waarmee het pyreen wordt
geïoniseerd dusdanig is, dat het ook op tijdschalen in de ruimte een belangrijke rol moet
spelen. Verder wordt in dit hoofdstuk getoond, dat het mogelijk is om kleine reactive
intermediairs in absorptie waar te nemen, die middels diffusie een belangrijke rol spelen
in de chemie van een PAK houdend ijs. Hoofdstuk 7 laat zien dat veel van de processen
zoals die voor pyreen in water ijs zijn gevonden, ook plaats vinden voor andere PAKs:
anthracene, coronene en benzo[ghi]perylene. De belangrijkste conclusie is, dat PAK chemie niet a priori buiten beschouwing mag worden gelaten in astrochemische modellen en
wellicht een belangrijke rol speelt bij de vorming van complexe organische moleculen.
Tenslotte bespreekt hoofdstuk 8 het perspectief van het beschreven onderzoek.
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Publications
Refereed papers
• Band profiles and band strengths in mixed H2 O:CO ices
J. Bouwman, W. Ludwig, Z. Awad, K. I. Öberg, G. W. Fuchs, E. F. van Dishoeck
and H. Linnartz
Astronomy & Astrophysics, 476, 995-1003 (2007) (Chapter 2)
• The c2d Spitzer spectroscopic survey of ices around low-mass young stellar objects.
IV. NH3 and CH3 OH
S. Bottinelli, A. C. A Boogert, J. Bouwman, M. Beckwith, E. F. van Dishoeck, K I.
Öberg, K. M. Pontoppidan, H. Linnartz, G. A. Blake, N. J. Evans II and F. Lahuis
Astrophysical Journal, 718, 1100-1117 (2010) (Chapter 3)
• Photochemistry of PAHs in cosmic water ice. Part I: Mid-IR spectroscopy and
photoproducts
J. Bouwman, A. L. Mattioda, H. Linnartz and L. J. Allamandola
Astronomy & Astrophysics, submitted (2010) (Chapter 4)
• Optical spectroscopy of VUV irradiated pyrene:H2 O ice
J. Bouwman, D. M. Paardekooper, H. M. Cuppen, H. Linnartz and L. J. Allamandola
Astrophysical Journal, 700, 56-62 (2009) (Chapter 5)
• Photochemistry of the PAH pyrene in water ice: the case for ion-mediated solidstate astrochemistry
J. Bouwman, H. M. Cuppen, A. Bakker, L. J. Allamandola and H. Linnartz
Astronomy & Astrophysics, 511, A33+ (2010) (Chapter 6)
• Photochemistry of PAHs in cosmic water ice. Part II: Near UV/VIS spectroscopy
and ionization rates
J. Bouwman, H. M. Cuppen, L. J. Allamandola and H. Linnartz
Astronomy & Astrophysics, in prep. (2010) (Chapter 7)
179
Publications
• High-resolution infrared spectroscopy of the charge-transfer complex
[Ar − N2 ]+ : A combined experimental theoretical study
H. Verbraak, J. N. P. van Stralen, J. Bouwman, J. S. de Klerk, D. Verdes,
H. Linnartz and F. M. Bickelhaupt
Journal of Chemical Physics, 123, 144305 (2005)
Conference Proceeding
• VUV photochemistry of PAHs trapped in interstellar H2 O ice
J. Bouwman, H. M. Cuppen, L. J. Allamandola, H. Linnartz, PAHs and the Universe, 2010
Popular article
• Zwerven door het onderwijs
J. Bouwman, Nederlands Tijdschrift voor Natuurkunde, November 2006
180
Curriculum vitae
I was born on May 3, 1979 in the city of Haarlem, the Netherlands. In the year 1991, I
started high school at het Mendelcollege. Here, I finished the MAVO in 1995.
I continued my education at the IJmond Middelbare Technische School in SantpoortNoord, where I did a technical study dedicated to industrial processing techniques. I
started this study in 1995 and after 3 theoretical years, I concluded this study with two
half-year long traineeships at Shell Research and Technology Center Amsterdam and at
Sigma Coatings Amsterdam. The research aspect of the internship at Shell was very
appealing to me and made me decide to pick up a new study rather than starting a career
in the industry.
After finishing the IJmond MTS in 1999 I chose to continue at the Technische Hogeschool Rijswijk. Here, a study at the level of Bachelor of Engineering called Technical
Physics was offered, which aimed to using physics to solve industrial problems. I chose
the specialization Applied Physics and did a one year internship at the Laser Centre Vrije
Universiteit. The internship took place at the Department of Physical Chemistry, led by
prof. dr. S. Stolte and my daily supervisor was dr. Harold Linnartz. The project consisted
of building an electron impact plasma setup, combined with a sensitive high resolution
infrared spectrometer with the goal to measure infrared spectra of van der Waals bound
ionic complexes. During this project I worked closely together with dr. H. E. Verbraak
and dr. J. S. de Klerk. I finished my degree with honors in 2004 on a bachelor thesis
entitled “High Resolution Infrared Absorption Spectroscopy of Ionic Complexes”. This
work resulted in a publication in a scientific journal.
After finishing my Bachelor degree on a research project at the Vrije Universiteit
Amsterdam, the obvious choice was to continue with a Master study. I started the study
Chemistry with specialization in Laser Sciences in February 2004. My major research
project took place at the Laser Centre, again under supervision of dr. Harold Linnartz.
It involved the construction of an experimental setup for performing cavity ring-down
spectroscopy on a planar plasma expansion for the spectroscopic detection of diffuse interstellar band carriers. During this project I worked closely together with dr. D. Ityaksov.
I graduated on my thesis entitled “High Resolution Electronic Spectroscopy of Molecular
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Curriculum vitae
Transients of Astrophysical Interest” in February 2006.
Now, having my Masters degree, I was offered a PhD project in the Raymond & Beverly Sackler Laboratory for Astrophysics, under supervision of again prof. dr. Harold
Linnartz. The project comprised the design and construction of a new setup for performing spectroscopy on vacuum ultraviolet processed interstellar ice analogues. This
project was financed both by FOM and NOVA. I studied the physical interactions and
chemical reactions in interstellar ice analogues, using spectroscopic techniques covering
the near ultraviolet to the mid infrared. I worked closely together with dr. L. J. Allamandola from NASA Ames Research Center; a collaboration which led to a two months
research visit at NASA Ames. The research focus in this period was on PAH containing
ices. During my PhD I guided several bachelor and master students and presented my
results at (inter)national scientific meetings and during colloquia in Amsterdam, Groningen, Leiden, Lunteren, Rolduc (the Netherlands), Toulouse (France), Columbus (OH),
Berkeley (CA), and NASA Ames Research Center (CA).
November 1st 2010, I will start a postdoctoral research project in prof. dr. S. R.
Leone’s group at UC Berkeley, where I will perform gas phase reaction rate and isomer
specific reaction product branching ratio measurements.
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Nawoord
Een proefschrift kan alleen tot stand komen met steun van anderen. Het was een voorrecht om aan de Sterrewacht Leiden te werken met mensen met een grote verscheidenheid
aan nationaliteiten en vakdisciplines. Ten eerste wil ik graag de ondersteunde kracht van
de computergroep en het secretariaat noemen. Het bouwen van de opstelling was nooit
gelukt zonder de technische ondersteuning van de fijnmechanische werkplaats. Ewie,
Gijsbert en Martijn, het was geweldig om samen met jullie de opstelling operationeel te
krijgen. De elektronische werkplaats, onder leiding van René, heeft ook zijn aandeel gehad in de het opbouwen en werkend krijgen van de opstelling. Additionally, I am grateful
to Andy Ruth for having me in his laboratory at University College Cork. The support
from, and interaction with, colleagues from the laboratory and astrochemistry group have
played a crucial role in my development, both personally and professionally. Sergio, you
have been a great colleague and friend during 4 years of my PhD. Harald, het was fijn
om met jou samen te werken, zowel aan de VU in Amsterdam als in het laboratorium
in Leiden. Working in the office has been great because of my three roommates Rafael,
David and Emily. The other members of the Sackler Laboratory group, Nadine, Edith,
Karoliina, and Joseph and the Sackler Lab alumni Guido, Suzanne, Karin and Claire have
also been of great support. Zonder Herma’s kennis van het modelleren van chemische
reacties hadden mijn publicaties lang niet een dergelijk grote impact gehad. Naast de
medewerkers van de universiteit heeft ook een aantal studenten — Wiebke, Martha, Daniël, Nienke, Arthur en Michiel — met mij samengewerkt aan het assembleren van de
opstelling, reduceren van data, schrijven van software en/of het uitvoeren van metingen.
Sandrine, working with you on the ammonia paper was a pleasure. Ook Ewine heeft een
zeer belangrijke rol gespeeld bij het voltooien van mijn proefschrift en het vergroten van
mijn kennis in de astrochemie. Xander, de gesprekken die wij hebben gehad, hebben een
belangrijke bijdrage geleverd aan de uiteindelijke keuze die ik heb gemaakt. Lou, it was
an honor and a pleasure to work with you in Leiden. I will never forget our Kasteel afternoons encompassing a perfectly balanced mixture of personal and work-related conversations. Voordat ik als promovendus naar de Universiteit Leiden kwam, heb ik bijzonder
veel vertrouwen en steun mogen ontvangen van Harold. Dit heeft er mede toe geleid dat
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Nawoord
nu dit proefschrift gedrukt voor u ligt. Furthermore, I like to wish my successors, Emily
and Steven, the best of luck with the beautiful and well-behaved system!
The two months visit to the astrochemistry group at NASA Ames Research Center
has been a life changing experience. The research I have done there seamlessly connects
part I and II of this thesis. Lou, Andy and all other (lab) colleagues, thank you for your
hospitality, for helping me with the measurements, and for taking me on the great trips
through beautiful CA. Additionally, I mention Christiaan and Claire, who have been kind
enough to introduce me to Mountain Views night life.
Naast collega’s en familie hebben ook mensen uit mijn sociale kring hun bijdrage
geleverd. Om te beginnen hebben de spelers van het zaalvoetbalteam van ‘de S’ mij de
nodige inspanning én ontspanning bezorgd. Helaas verdwenen uit het betaald voetbal,
maar niet uit mijn hart, HFC Haarlem, en de vele vrienden die ik daar aan over heb
gehouden. Het was fantastisch om tweewekelijks onze passie te delen en ik zal nooit
onze reisjes door Nederland en daar buiten vergeten.
Graag noem ik mijn vriend en medepromovendus Tom, waar ik bijzonder veel steun
aan heb gehad. Het delen van onze passies in het privé leven, gecombineerd met het
delen van onze ervaringen in de wetenschap hebben er mede voor gezorgd dat dit mooie
proefschrift tot stand is gekomen.
Tot slot wil ik mijn familie noemen, waaronder een aantal mensen in het bijzonder.
Marc en Anita, ondanks onze drukke schema’s is het ons altijd gelukt tijd vrij te maken
voor gezellige avonden! Pap, mam, Jeroen en Suus, die mij altijd mijn eigen keuzes
hebben laten maken en mij hebben gesteund door dik en dun! De familie van mijn lieve
vrouw, Klaas, Ingrid, Jeroen, Marc en Remco, die mij met hun nuchtere kijk op zaken
ook op de rails hebben gehouden. Mijn proefschrift was nooit tot stand gekomen zonder
de rust, liefde en steun van mijn lieve vrouw Wendy. Je hebt ervoor gezorgd dat ik ook in
drukke tijden mijn focus behield en mijn rustmomenten pakte!
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