Lectures 7-8 ISM: Gas

Lectures 7-8 ISM: Gas
GAS
The other baryonic component- sec 2.4 in
the text
log gas fraction for star
forming galaxies
Gas
• Other than stars the baryons in
galaxies lie in 3 forms
– gas
– rocks
– dust (0.1% of mass)
• the % mass in rocks and dust is
small
• There is an interplay between the
stars and gas,with stars forming out
of the gas and with enriched gas
being ejected back into the
interstellar medium from evolved
stars.
• There exist a vast array of spectral
diagnostics for the gas in both
emission and absorption which can
reveal
– chemical composition
– temperature
– velocities
– ionization mechanism
Peeples and Shankar 2011
Gas- Big Picture
• dark matter halos grow by merging and accretion (e.g.
Galaxies can grow by accretion of gas, by merging with gas
rich galaxies and by merging with gas poor galaxies)
• gas falls into these halos, cools and forms stars.
• How does this occur- the physics of gas accretion,
– How and when did galaxies accrete their gas and what
do they do with it (e.g. form ISM, stars, expel the gas,
feed the supermassive black hole ....)
• ‘cold’ gas: dominates in Spiralsmany phases
– neutral hydrogen
– molecular gas-Dense molecular
clouds, have most of the total
mass of the interstellar gas and
are of key importance for star
formation, occupy a negligible
fraction of the total volume
– warm ionized gas-has persistent
transient states out of thermal
pressure balance
• Milky-Way-like galaxies cold gas
mass~10% of the stars
• For lower mass galaxies the
baryonic fraction in gas is larger; at
Mhalo<1011M
gas dominates the baryonic content
• Hot gas (T~106-7 k) dominant ISM
in elliptical galaxies
In spirals hot gas fills the volume but
low total mass
GAS-ISM
x-ray images of elliptical galaxies
emphasing structure
Physics of Emission from Gas
• Gas is heated/excited/ionized by photons (stars, AGN), shocks
(supernova)
• Atomic transitions reveal the ionization state, temperature, density,
velocity structure and chemical composition of the gas.
• Photoionization: photon from source eject electron from ion- to do
this photon needs to have energy greater than ionization potential
(e.g. 13.6 eV for Hydrogen; O,B stars, AGN)
• Collisional ionization: gas is excited by collisions with 'hot' electrons
(again electron energy has to be above threshold). Electrons have
Maxwell-Boltzman energy distribution in equilibrium
• wide range of types of transitions: 2 'basic' types
– permitted: fast transition rate, line is emitted before ions state is
altered
– forbidden: violate transition rule, ion can be collisionally deexcited when density exceeds critical density; presence of line
thus places constraint on gas density. - jargon forbidden lines are
indicated by [OII] (OII is the ionization state of the gas, once
ionized oxygen).
A Bit of Physics
• Saha equation describes the
ionization balance of the gas which
depends on the temperature,
quantum mechanical transition
probabilities and densities
• An atom with multiple energy states
in thermal equilibrium with a
radiation field will find itself in one or
another of these energy states.
• Frequent transitions to and from
other states will occur as photons
interact with the atoms.
• transitions from the upper of the
states of figure take place by photo
deexcitation and by induced
deexcitation. Transition in the
upward direction is by
photoexcitation or collisional
excitation
For lots of details
see MBW appendix B
A Bit of Physics
• The rates of ionization and
recombination are important (see eqs
2.21,2..22 in S+G); e.g. X++e
X+γ
• the rate at which ions recombine thus
clearly depends on the ion density .
X+and the electron density and the
recombination coefficient α, which
depends on the ion, (e.g. the number
of electrons it has and its atomic
number)
• Thus recombination rate of electrons
for a given ion X++ is
• dne/dt=nx+neα(Te);
• the recombination time is the #of
electrons/ the rate :ne/dne/dt
a few thousand years in a HII region α the recombination rate depends on
QM and Boltzmann's law
In steady state # of
ionizations= # of
recombinations
Ionization is from
• collisions with hot electrons
•photoionization from stars
•shocks
Atomic Lines
• The energy levels and
transitions for hydrogen
• Each element and ionization set
has a similar (but more complex)
set of lines
• The probability of emitting a
given line depends on the
temperature and density of the
gas
A Bit of Physics-Ionizing Photons
• One can estimate the number of
ionizing photons from a star
using the black body formula
(e.q. 1.35 in text) and
integrating over the photons
more energetic than the
ionization potential of the ion of
interest (e.g. H with13.6 eV)
• These photons ionize and heat
the gas
• The gas responds by emitting
lines characteristic of the
chemical composition,
temperature, ionization state,
density etc ...
Importance of the ISM
• Despite its low mass, the ISM is
very important
• crucial role in the star-gas cycle
in spirals and irregulars,
– it facilitates ongoing (&
current) star formation
– it is a repository for
elements created in SNR and
stars and therefore is a key
to measure chemical
evolution
• Because it can cool, its collapse
is dissipational
– stars can form !! hot gas
cold gas
stars:
– galaxies are smaller than
dark matter halos !
its emission & absorption provides
enormous diagnostic information
•Doppler motions reveal galaxy
dynamics
•Abundance measurements allow
study of chemical evolution
•physical conditions: density; temp;
pressure; turbulence; gas column
density; mass,
•can all be derived from
observations of emission
lines
•lines are bright and can be
seen (relatively) easily at
cosmological distances.
The ISM in Spirals is DYNAMIC
•
•
There is strong
interaction
between the
different phases of
the ISM and
feedback between
star formation and
the rest of the ISM
There is lots of
complex nonlinear effects (and
lots of jargon)
Fabian Walter
At low redshift ISM in
spirals not affected much by AGN
Its not so clear if ISM in
ellipticals is dynamic in the
same way; AGN seem to be more important
Big Questions
• What is the volume filling factor
of the hot ISM?
• What is the distribution of the
temperature, density, and
velocity
• What are typical scales in the
ISM and why?
• What is the effect of turbulence ,
magnetic fields and cosmic rays
• What causes density and
pressure inhomogeneities in the
evolution of the ISM
• How is the ISM related to star
formation
temperature
map of LMC
intensity
map of LMC in
x-rays
Spiral ISM 'States'- f is the filling factor
• Molecular Medium (MM): T~20 K, n > 103 cm-3, f < 1%. The MM is
mostly cold dense molecular clouds which are gravitationally bound.
Although, on average, this phase contains as much mass as the atomic
hydrogen, it occupies only a very small fraction of the ISM.
• Cold Neutral Medium (CNM; T~100 K, n~20 cm-3, f =2 - 4%). The CNM
is distributed in rather dense filaments or sheets, occupying a minor
fraction of the ISM. The CNM is most readily traced by HI measured in
absorption.
• - Warm Neutral Medium (WNM; T~6000 K, n~ 0.3 cm-3, f~30%). This
phase provides the bulk of the HI seen in emission line surveys.
• - Warm Ionized Medium (WIM; T ~8000 K, n ~0.3 cm-3 f~15%).
associated with HII regions, but a considerable fraction of the ISM
outside of HII regions is also filled with ionized gas.
• - Hot Ionized Medium (HIM; T~106 K, n~10-3 cm-3 , f~50%). The hot gas
produced by supernova explosions and their after effects - long cooling
time , a large fraction of the ISM is filled with this component.
• http://ned.ipac.caltech.edu/level5/March01/Brinks/Brinks4.html
ISM- Phases
• Hot ionized medium (e.g. X-rays)
• Warm ionized medium HII region(e.g. Hα)
• Warm neutral medium (e.g. HI emission)
• Cold neutral medium (e.g. HI absorbtion)
• Molecular medium (e.g. CO)
These phases have different distributions perpendicular to the planescale height
Fabian Walter
The ISM
•
•
•
•
The 5 'states' are in dynamic interaction.
the coldest clouds are molecular and the densest (hydrogen molecules, CO, NH3 and other
molecule)s- this is where stars form .
The dust is composed of 'refractory' elements and molecules mainly carbon, silicon, iron
and is responsible for most of the absorption of optical light in the galactic plane - the
energy absorbed by the dust heats it and the dust re-radiates in the IR
The ISM is threaded by magnetic fields. At ~ 5µG, these fields provide a pressure
comparable to the pressure of the gas . The magnetic fields therefore affects the dynamics
of the ISM
• How dense (or not) is the ISM?
– Dense molecular cores: n~105 cm-3
– Typical ISM: n~1cm-3
– sea level earth atmosphere: 1 mol / 22.4L -3x1019cm-3 (oxygen)
– Best human-made vacuum:3x104cm-3
Optical spectrum of HII Region
• Optical spectrum show lines due to [OII]. [OIII],Hα, [NII], etc
Millimeter Band Spectrum of Molecular
Cloud
Millimeter Band Spectrum of Molecular
Cloud
• zoom in on previous plot
Molecular Lines
• Molecular clouds are very rich in spectral features from a wide
variety of molecules- lots of information
• Some of the lines (CO) are so strong that they can be seen at high
redshift
ISM- Relevant Velocities
Some characteristic values
• galactic rotation speed 18km/sec/kpc
• Thermal sound speed ideal gas cs=sqrt(kT/mHµ); µ is mean molecular
weight
for Hydrogen this is 1km/sec at 100k
Alfven speed- speed at which magnetic fluctuations propagate
vA=B/sqrt(4πρ)- so for typical ISM values B=1µG, ρ=1cm-3
vA= 2 km/sec
Physics of Emission from Gas
• Lines have enormous range of energies/wavelengths
– molecular and fine structure lines in IR/radio band
– atomic lines in the IR, optical, UV and x-ray
• Ionized gas also emits a continuum via thermal bremmstrahlungshape of which is a measure of temperature, intensity goes as density
squared.
• Observed line energies give velocity information: redshift, velocity
field
• Relative strength of lines determines ionization temperature,
abundance of given element (corrected for ionization balance (go to
board)).
• see Space Science Reviews, Volume 134, Issue 1-4, pp. 155-190,
2008m astro-ph/0801.1011 for the background physics for low
density, hot gas.
–
Gas Cooling L=n2Λ(Τ)
• T>107k thermal
bremmstrahlung L~n2T 1/2
• 107>k T>106.3k Fe L lines
• 104.5>kT>106.3k K and L lines of
'metals'
• 104>kT>104.5k
• At lower temperatures fine
structure lines and molecules
take over
log T
Cooling curve as a function
of kT and metallicity-for gas
in collisional equilibrium
Sutherland and Dopita
table 2.5 in text
Gas Cooling
As the temperature changes
the ions responsible for cooling
change as do the physical
processes
http://www.cfa.harvard.edu/swas/swasscienc
e/fig2.html
Gas Cooling
• Collisional excitation: free
electron impact knocks a bound
electron to an excited state; it
decays,emitting a photon.
• Collisional ionization: free
electron impact ionizes a
formerly bound electron, taking
energy from the free electron.
• Recombination: free electron
recombines with an ion; the
binding energy and the free
electron's kinetic energy are
radiated away
• Free-free emission: free electron
is accelerated by an ion,
emitting a photon. (A.k.a.
Bremsstrahlung.)
Thermal Bremmstrahlung
• Electrons have a MaxwellBoltzmann distribution
• electromagnetic radiation
produced by the deceleration of
a charged particle when
deflected by another charged
particle, typically an electron by
an atomic nucleus (wikipedia)
• Bremsstrahlung has a
continuous spectrum, whose
shape depends on temperature
roughly E-0.4 exp(-E/kT)
Gas Cooling
• The functions are very different
for photoionized gas not in
collisional equilibrium
• This depends on the shape of the
photon spectrum and its
intensity
• This is very important for
studies of active galaxies and
the intergalactic medium
Gas Heating Mechanisms in ISM
• heating by low-energy cosmic
rays (dense MM)
• photoelectric heating by grains
(CNM to MM)
• photoelectric heating by
photoionization of atoms and
molecules (HII regions)
• photoelectric heating by soft Xrays (WIM, WNM, CNM)
• chemical heating (dense MM)
• grain-gas thermal exchange
(dense MM)
• hydrodynamic and
magnetohydrodynamic heating
(WNM, CNM)
• interstellar shocks (WNM, CNM,
MM) due to supernova
X-ray image of Cas-A
youngest SNR in MW
Fabian Walter
in H α
Cosmic Rays
•
Cosmic rays, which are atomic nuclei electrons and protons which have been
accelerated to nearly the speed of light- thought to be created in SNR shocks
• Gyroradius=rg=p/qB
(p is the momentum of the particle, B the magnetic field, q the charge)
In handier units r=3.3x107γ/B(gauss)cm ; γ is the relativistic factor sqrt(1/(1-v2/c2))
With B~5uG the gyroradius of a proton with γ~104 (a typical value) is ~10-4 pc.
so cosmic rays are trapped within the Galaxy by the magnetic fields .
Despite the wide range of densities and temperatures in the ISM, the pressure in the
plane is roughly uniform .
The pressure in the ISM counteracts the gravitational force of the Galactic disk . It is
highest in the plane and drops as higher galactic latitudes - but the scale height of
the drop is related to the temperature of the gas.
The ISM can dominate a galaxy's integrated SED -in the far IR and radio
• Mid-IR to Sub-mm is dominated by emission from dust, molecular lines
and fine structure lines
• radio comes either from HII regions or a relativistic plasma radiating
via synchrotron radiation
certain emission lines (eg Lyα ; [CII] 158µ) can be major coolants
ISM in Spirals
• The ISM is energized primarily by
stars (starlight (dust), stellar winds,
supernovae,
• UV starlight photoionizes atoms &
dissociates molecules; photo-ejected
electrons heat gas
• SN shocks heat/ionize/accelerate
gas & are largely responsible for the
ISM's complexity in spirals.
• The interstellar medium near the
Sun has large scale structures of
bubble walls, sheets, and filaments
of warm gas.
• The remainder of the volume is in
bubble interiors, cavities, and
tunnels of much lower density, hot
enough to be observable via their Xray emission (Cox ARA&A)
ISM in Spirals
• The phases of the gas are
distributed differently
– cold (molecular) gas is
confined to a thin disk
ρ(z)~ 0.58 exp[-(z/81 pc)2]
and has a mean T~15k
– 'warm' gas has a density
distribution
total gas density in MW
vs height above the
disk
(blue)
warm gas in red
ρ(z)~ 0.57 * 0.18 exp[-(z/318 pc)2]
where z is the distance above
the disk midplane
has a mean T~5000k
Roughly magnetic (~5µG),
cosmic ray, and dynamical
pressures are equal ~10-12
dyne mid-plane
Cox Ann Rev A&A
X-ray ISM in M101
• Hot phase of ISM in M101dominated by ionized oxygen
OVII/OVIII and T~2x106k is the
temperature of the dominant
component.
• The emission is centrally
concentrated
• Such data exists for only a few
objects
x-ray
surface
brightness
x-ray image
ISM In Ellipicals-pg
272 in S+G
• Predominately hot kT~106-107K
and thus visible only in the xray
– the temperature is set,
predominantly by the depth
of the potential well of the
galaxy (if it were hotter it
would escape, if colder fall)
– The metallicity of the gas is
roughly solar
x-ray image of NGC 4636
X-ray Spectra of NGC1399
• At certain temperatures (~416x106k) the spectrum is
dominated by Fe lines from the
L shell whose energy is very
sensitive to temperature.
• Thus x-ray images and spectra
(obtained simultaneously with
CCDs) get the density and
temperature and estimates of
the chemical composition of the
ISM in ellipticals
wavelength Å
Hot Gas and Metallicity
• In elliptical galaxies ,
clusters of galaxies and
star forming galaxies the
ISM is hot and emits
primarily via thermal
bremmstrahlung with
strong emission lines from
abundant elements (O, Ne,
Si, S, Fe)
• These are fairly easy to
measure and the amount
of hydrogen is measured
by the strength in the
continuum.
• Problem is x-ray sources
are weak and telescopes
are small so not so many
objects (~100's)
X-ray spectrum of hot
gas in a star forming galaxy
Image of x-ray source
determines
the gas density since L~Λ(Τ)n2
Metallicity in Gas
• For star forming galaxies it is easier to
measure the metallicity in the gas
phase than in the stars-strong
emission lines-but one measures
different elements
• How does one do it ?- Use HII region
spectra (ionized gas around hot young
stars) : measure oxygen lines.
• O is an α-process element made in
short-lived massive stars and is~50%
of all the heavy elements by mass representative of all the heavy
elements made in type II SN
– need to measure line strengths,
electron temperature, density to
get ionization structure of the gas
(see Lopez-Sanchez et al 2012)
–
• More massive galaxies tend to be more
metal rich
There are several methods
to do this- but error of factor
of 2
MW
log Massstars
GAMA collaboration Foster et
al 2012
Metallicity Issues for Distant Galaxies
Fundamental problem is that the ionization structure of the gas is
unknown and the line strengths and hence the abundances depend on
both
• chemical abundances,
• the ionization parameter ( U) which is the ratio of ionizing photons
density to gas density for photoionized gas:(U~L/ner2)
log [OIII}/[OII]
each line corresponds to the
predicted [OIII]/[OII] ratio for a
different abundance (0.05-3x
solar) and ionization parameter
A fixed line ratio can correspond
to a factor of 20 range in
abundance if ionization parameter
is not simultaneously constrained.
ionization parameter
Metallicity Issues for Distant Galaxies
Since the electron temperature ,density and nature of stellar ionization
field vary quite a bit over the galaxy these are 'irreducible' errors.
One resorts to calibrating the lower quality galaxy data against the
excellent data for HII regions in the MW and some other nearby
galaxies
Gas phase abundances are 'ok' for O,N and S (but not Fe)
Abundances determined in stars mainly measure 'Fe' via absorption lines
in stellar spectra (Worthy et al 1994)- very very messy.
Why Metals are Important
• metals account for 1% of the
mass, they dominate most of
the important chemistry,
ionization, and heating/cooling
processes.
• Comparison of the metal content
of gas and stars compared to
– what is expected from stellar
evolution
– cosmic star formation rates
indicates whether galaxies
expel metals and/or accrete
gas.
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