76 T dwarfs from the UKIDSS LAS: benchmarks, kinematics

76 T dwarfs from the UKIDSS LAS: benchmarks, kinematics
MNRAS 433, 457–497 (2013)
doi:10.1093/mnras/stt740
Advance Access publication 2013 June 4
76 T dwarfs from the UKIDSS LAS: benchmarks, kinematics
and an updated space density
1 Centre
for Astrophysics Research, Science and Technology Research Institute, University of Hertfordshire, Hatfield AL10 9AB, UK
Nacional, Rua General José Cristino, 77 – São Cristóvão, Rio de Janeiro – RJ 20921-400, Brazil
3 Istituto Nazionale di Astrofisica, Osservatorio Astrofisico di Torino, Strada Osservatorio 20, I-10025 Pino Torinese, Italy
4 Gemini Observatory, 670 N. A’ohoku Place, Hilo, HI 96720, USA
5 Institute for Astronomy, University of Hawai‘i, 2680 Woodlawn Drive, Honolulu, HI 96822, USA
6 Department of Astronomy, Boston University, 725 Commonwealth Ave, Boston, MA 02215, USA
7 Australian Centre for Astrobiology, University of New South Wales, NSW 2052, Australia
8 School of Physics, University of New South Wales, NSW 2052, Australia
9 Department of Astronomy and Astrophysics, University of California, Santa Cruz, CA 95064, USA
10 Los Alamos National Laboratory, PO Box 1663, MS F663, Los Alamos, NM 87545, USA
11 Department of Physics and Astronomy, University of Sheffield, Sheffield S3 7RH, UK
12 Institut fur Astrophysik, Georg-August-Universitat, Friedrich-Hund-Platz 1, D-37077 Gottingen, Germany
13 C.R.A.L. (UMR 5574 CNRS), Ecole Normale Superieure, F-69364 Lyon Cedex 07, France
14 Instituto de Astrofı́sica de Canarias (IAC), Calle Vı́a Láctea s/n, E-38200 La Laguna, Tenerife, Spain
15 Instituto de Astrofı́sica de Canarias, E-38200 La Laguna, Spain
16 Max Planck Institute for Astronomy, Koenigstuhl 17, D-69117 Heidelberg, Germany
17 NASA Ames Research Center, Mail Stop 245-3, Moffett Field, CA 94035, USA
18 Shanghai Astronomical Observatory/CAS, 80 Nandan Road, Shanghia 200030, China
19 Observatório do Valongo/UFRJ, Ladeira Pedro Antonio 43, Rio de Janeiro – RJ 20080-090, Brazil
20 Universidad de Chile, Camino el Observatorio # 1515, Santiago, Casilla 36-D, Chile
21 Harvard–Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
22 Subaru Telescope, 650 North A’ohoku Place, Hilo, HI 96720, USA
23 Fundación Galileo Galilei – INAF, Apartado 565, E-38700 Santa Cruz de La Palma, Spain
24 Department of Astrophysics, American Museum of Natural History, Central Park West at 79th Street, New York, NY 10024, USA
25 National Astronomical Observatory, Mitaka, Tokyo 181-8588, Japan
2 Observatório
Accepted 2013 April 26. Received 2013 April 23; in original form 2013 February 7
ABSTRACT
We report the discovery of 76 new T dwarfs from the UKIRT Infrared Deep Sky Survey
(UKIDSS) Large Area Survey (LAS). Near-infrared broad- and narrow-band photometry and
spectroscopy are presented for the new objects, along with Wide-field Infrared Survey Explorer
(WISE) and warm-Spitzer photometry. Proper motions for 128 UKIDSS T dwarfs are presented
from a new two epoch LAS proper motion catalogue. We use these motions to identify two
new benchmark systems: LHS 6176AB, a T8p+M4 pair and HD 118865AB, a T5.5+F8 pair.
Using age constraints from the primaries and evolutionary models to constrain the radii, we
have estimated their physical properties from their bolometric luminosity. We compare the
colours and properties of known benchmark T dwarfs to the latest model atmospheres and
draw two principal conclusions. First, it appears that the H − [4.5] and J − W2 colours are
more sensitive to metallicity than has previously been recognized, such that differences in
E-mail: [email protected]
C 2013 The Authors
Published by Oxford University Press on behalf of the Royal Astronomical Society
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Ben Burningham,1,2‹ C. V. Cardoso,1,3 L. Smith,1 S. K. Leggett,4 R. L. Smart,3
A. W. Mann,5 S. Dhital,6 P. W. Lucas,1 C. G. Tinney,7,8 D. J. Pinfield,1 Z. Zhang,1
C. Morley,9 D. Saumon,10 K. Aller,5 S. P. Littlefair,11 D. Homeier,12,13
N. Lodieu,14,15 N. Deacon,16 M. S. Marley,17 L. van Spaandonk,1 D. Baker,1
F. Allard,13 A. H. Andrei,2,3,18,19 J. Canty,1 J. Clarke,1 A. C. Day-Jones,1,20
T. Dupuy,21 J. J. Fortney,9 J. Gomes,1 M. Ishii,22 H. R. A. Jones,1 M. Liu,5
A. Magazzú,23 F. Marocco,1 D. N. Murray,1 B. Rojas-Ayala24 and M. Tamura25
458
B. Burningham et al.
metallicity may dominate over differences in Teff when considering relative properties of cool
objects using these colours. Secondly, the previously noted apparent dominance of young
objects in the late-T dwarf sample is no longer apparent when using the new model grids and
the expanded sample of late-T dwarfs and benchmarks. This is supported by the apparently
similar distribution of late-T dwarfs and earlier type T dwarfs on reduced proper motion
diagrams that we present. Finally, we present updated space densities for the late-T dwarfs,
and compare our values to simulation predictions and those from WISE.
Key words: surveys – brown dwarfs – stars: low-mass.
1 I N T RO D U C T I O N
2 C A N D I DAT E S E L E C T I O N
Our initial candidate selection followed a similar method to that
described in Pinfield et al. (2008) and Burningham et al. (2010b),
which we summarize here. The selection process consists of two
channels: (1) those sources that are detected in the three UKIDSS
LAS YJH bands (the YJH channel); and (2) those sources that are
only detected in the YJ bands (the YJ-only channel). We did not employ WISE data for guiding our selection since the all-sky catalogue
became available part way through our follow-up campaign, and
consistent selection is crucial for establishing a well-characterized
statistical sample. Additionally, for fainter T6 and earlier dwarfs, the
WISE faint limits are effectively shallower than the UKIDSS LAS.
2.1 The YJH selection channel
Our YJH selection channel requires sources to lie within the
UKIDSS LAS sky that overlaps with the Sloan Digital Sky Survey (SDSS) DR8 footprint and have the following photometric
characteristics:
(i) J − H < 0.1
(ii) J − K < 0.1 or K band non-detection
(iii) z − J > 2.5 or no SDSS detection within 2 arcsec.
We also imposed a number of data quality constraints to minimize
contamination from artefacts and poor signal-to-noise (S/N) data
for which we refer the reader to Appendix A, which includes the
SQL queries we used to access the LAS via the WFCAM Science
Archive (WSA; Hambly et al. 2008). The epoch difference between
the UKIDSS data and SDSS data is variable and ranges up to six
years. Our inclusion of UKIDSS sources with no SDSS counterpart
within 2 arcsec introduced sources with z − J < 2.5 and proper
motions above ∼300 mas yr−1 to our candidate list. This source of
contamination was relatively small and in most cases such objects
were identified as fast moving earlier type objects prior to detailed
follow-up. Figs 1 and 2 show the UKIDSS LAS YJH photometry of
our selected candidates from DR5 to DR8. The greatest degree of
contamination in our YJH selection channel is from photometrically
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
The current generation of wide-field surveys is bringing about a
step change in our understanding of the coolest and lowest mass
components of the solar neighbourhood. The total number of cool T
dwarfs, substellar objects with 1400 Teff 500 K, has been taken
into the hundreds by infrared surveys such as the UKIRT Infrared
Deep Sky Survey (UKIDSS; Lawrence et al. 2007), the CanadaFrance Brown Dwarf Survey (CFBDS; e.g. Delorme et al. 2010)
and most recently the Wide-field Infrared Survey Explorer (WISE;
Wright et al. 2010). The last of these, which is an all-sky midinfrared survey, has extended the substellar census to well below
Teff = 500 K, and the adoption of a new spectral class ‘Y’ has been
suggested to classify these new extremely cool objects (Cushing
et al. 2011; Kirkpatrick et al. 2012). The VISTA Hemisphere Survey
(VHS; McMahon et al., in preparation) and VIKING survey are now
also adding to the census (Lodieu et al. 2012c; Pinfield et al. 2012).
Our exploitation of the UKIDSS Large Area Survey (LAS) has
focused on using the photometric characteristics of mid-to late-T
dwarfs at red optical and near-infrared (NIR) wavelengths (see e.g.
Kirkpatrick 2005 for a review of the L and T spectral classes) to select a statistically complete sample of T dwarfs across the T6–T8+
range (Lodieu et al. 2007b; Pinfield et al. 2008; Burningham et al.
2010b). This allowed us to identify an apparent dearth of late-T
dwarfs in the solar neighbourhood compared to Monte Carlo simulations based on functional forms of the initial mass function (IMF)
that have been fitted to observations of the substellar component of
young clusters (Pinfield et al. 2008; Burningham et al. 2010b).
Interpreting this result is hampered by the inherently indirect nature of the observations, and the problems associated with determining the properties of a mixed-age population of brown dwarfs, which
by their nature have no single mass–radius relationship. Leggett
et al. (2010a) found that colours of the late-T dwarfs that were
identified in UKIDSS suggest they are a predominantly young and
low-mass population from comparisons to the model atmospheres
of Saumon & Marley (2008). This surprising result could have significant bearing on the interpretation of observed space densities of
late-T dwarfs.
Benchmark brown dwarfs provide the opportunity to break the
degeneracies in age, mass and metallicity that hamper the characterization of cool substellar objects (Pinfield et al. 2006). As part
of our search of the UKIDSS LAS, and more recently WISE and
VISTA, we have identified several wide binary systems that allow
fiducial constraints to be placed on the properties of the T dwarf
secondary component (Burningham et al. 2009, 2011; Day-Jones
et al. 2011; Pinfield et al. 2012). Comparisons of these objects with
both the BT Settl (Allard, Homeier & Freytag 2010) model colours,
and those of Saumon & Marley (2008) also appear to support the
result of Leggett et al. (2010a).
In this paper, we present the results of the extension of our search
of the UKIDSS LAS up to and including the sky available in Data
Release 9 (DR9), with near-complete follow-up of Data Release 8
(DR8). We have used the significantly enhanced sample of late-T
dwarfs, along with two epoch UKIDSS LAS proper motions from
the catalogue of Smith et al. (in preparation) to perform a systematic search for wide binary benchmark objects, and compare the observed properties of the benchmark sample and the wider UKIDSS
sample with the latest model prediction from Saumon et al. (2012)
and Morley et al. (2012). We also use the proper motions for our
sample to briefly investigate the kinematics of UKIDSS T dwarfs.
Finally, we provide an updated estimate of the space densities of
T6–T8+ dwarfs, and discuss possible routes to reconciling the observations of the field and young clusters.
76 T dwarfs
459
scattered M dwarfs, with the greatest frequency of contaminants
found at faint J, blue Y − J and red J − H.
2.2 The YJ-only selection channel
Our YJ-only selection channel ensures that we do not exclude bona
fide late-T dwarfs that are fainter than the LAS H-band detection
limit due to the inherently blue J − H colours of such objects. To
minimize contamination from scattered and blue M dwarfs we impose a Y − J constraint for this selection, such that our photometric
criteria are:
(i) Y − J > 0.5 or J < 18.5
(ii) H and K band non-detection
(iii) z − J > 2.5 or no SDSS detection within 2 arcsec.
We applied the same data quality restrictions to the YJ-only channel as were applied to YJH selection channel. In addition, to minimize contamination from Solar system Objects (SSOs), which can
appear as non-detections in the H and K bands due to different
epochs of observation, we also imposed a criterion that Y- and Jband coordinates must agree to within 0.75 arcsec for observations
taken within a day of each other. For observations taken more than
a day apart, we remove this requirement to avoid excluding bona
fide candidates with high proper motion. We again refer the reader
to Appendix A for details of the SQL queries that we used to access
the WSA. The additional SSO contamination present in the YJ-only
selection channel can be seen in Fig. 2 as relatively bright and blue
Y − J contaminants with very blue limits on their J − H colours.
3 F O L L OW- U P P H OT O M E T RY
To remove contaminants such as photometrically scattered M
dwarfs and SSOs, we used a combination of NIR and optical pho-
Figure 2. JH colour–magnitude plot showing UKIDSS LAS photometry
of candidate T dwarfs from DR5 to DR8. Confirmed T dwarfs are shown
with blue filled circles, with arrows indicating limits on J − H colours for
candidates from the YJ-only channel. Limits are for illustrative purposes,
and are based on the canonical 5σ depth for the LAS of H = 18.8, this
results in these candidates forming a straight diagonal sequence across the
plot. Rejected candidates are shown with filled red squares, whilst yet-to-be
followed-up targets are shown with open circles. For reference, field stars
from a randomly selected 1 square degree region of LAS sky are shown as
black dots.
tometry. We followed two distinct strategies in removing contaminants. Prior to 2010 December, we continued the strategy previously described in Burningham et al. (2010b). Briefly, this involved
obtaining higher SNR H-band photometry to remove early-type objects that had been scattered into our J − H < 0.1 selection (or to fill
in the missing data for H-band drop-outs in the YJ-only channel),
and repeat J-band observations to remove SSOs. Whilst adequate
for drawing a roughly complete sample of T4+ dwarfs, this method
was unable to effectively prioritize objects with spectral types of
T6 and later for follow-up. Since the T6–T9 region is most useful
for constraining the form of the field mass function (e.g. Burgasser
2004), we revised our strategy to allow the rejection of most of the
T4 dwarfs that dominated our Burningham et al. (2010b) sample.
Our revised strategy involved using relatively short exposure (15–
30 min) z -band imaging to confirm red z − J > 2.5 colours for
targets with no SDSS detection and for which the limits were not
sufficient to rule out a bluer z − J colour. For targets from the YJonly channel, this step was preceded by short J-band observations
to remove SSO contaminants. Targets with confirmed z − J > 2.5
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 1. YJH colour–colour plot showing UKIDSS LAS photometry of
candidate T dwarfs from DR5 to DR8. Confirmed T dwarfs are shown with
blue filled circles, with arrows indicating limits on J − H colours for candidates from the YJ-only channel. Limits are for illustrative purposes, and are
based on the canonical 5σ depth for the LAS of H = 18.8. Rejected candidates are shown with filled red squares, whilst yet-to-be followed-up targets
are shown with black open circles. For reference, field stars from a randomly selected 1 square degree region of LAS sky are shown as black dots.
460
B. Burningham et al.
were then targeted with CH4 imaging to identify late-type T dwarfs.
Targets with CH4 s − CH4 l < −0.5 were prioritized for spectroscopy
based on the methane colours of T dwarfs reported by Tinney et al.
(2005).
In the following subsections, we outline the photometric observations and data reduction that were carried out for this follow-up
programme. Details of the observations carried out for each target
are given in Appendix B.
3.1 Broad-band photometry
3.2 CH4 photometry
Differential methane photometry was obtained using the Near Infrared Camera Spectrometer (NICS; Baffa et al. 2001) mounted on
the TNG under programme AOT22 TAC 96 spanning from 2010
to 2012. NICS contains a set of MKO NIR filters as specified by
Tokunaga et al. (2002). More information about the narrow-band
and intermediate-band sets are given by A. Tokunaga.2 The methane
filters used in this work are denoted as CH4 s and CH4 l. The comparison of these two filters provides information about the strength
of the methane absorption bands in late-T dwarfs. CH4 l samples the
methane absorption bands present between 1.6 and 1.8 µm, while
the CH4 s samples a pseudo-continuum outside the methane band.
The final image mosaics were produced using the Speedy Nearinfrared data Automatic Pipeline (SNAP) provided by TNG (version
1.3). SNAP is an automated wrapper of existing pieces of software
(IRDR, IRAF, SEXTRACTOR and DRIZZLE) to perform a full reduction
with a single command. SNAP performs flat-fielding, computes the
offsets between the dithered images, creates a mosaic image with
double-pass sky subtraction and correction for field distortion.
The data were taken in different observing conditions, from photometric conditions to cirrus. Photometry was performed with IMCORE, part of CASUTOOLS (version 1.0.21), using a fixed circular aperture of 2 arcsec. CASUTOOLS3 is a suite of programs developed and
used by the CASU for survey data reduction tasks associated with
the UKIDSS and VISTA surveys, amongst others.
Differential photometric calibration of the methane colour
CH4 s − CH4 l was performed using the UKIDSS field stars present
in the field, and the method defined by Tinney et al. (2005) in their
section 2.5. Tinney et al. (2005) only provides the parametrization
for the 2MASS system. Using the information available on their
table 3, we performed the parametrization for the UKIDSS system
avoiding the region 0.48 < (J − H)MKO < 0.512 where the sequence
is degenerate. The sequence was fitted with two separate quadratics
to the regions −0.050 < (J − H)MKO < 0.480,
CH4 s − CH4 l = +0.000 46 − 0.012 59(J − H )
+ 0.318 17(J − H )2
and 0.512 < (J − H)MKO < 1.000,
CH4 s − CH4 l = −0.173 17 + 0.927 44(J − H )
− 0.589 69(J − H )2 .
2
1
http://www.oracdr.org/
(1)
3
(2)
See http://www.ifa.hawaii.edu/∼tokunaga/NB_special_ordersorting.html.
http://apm49.ast.cam.ac.uk/surveys-projects/software-release
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Our broad-band NIR photometry was obtained using the UKIRT
Fast Track Imager (UFTI; Roche et al. 2003) and WFCAM (Casali
et al. 2007), both mounted on UKIRT across a number of observing runs spanning 2009 to the end of 2010. UFTI data were dark
subtracted, flat-field corrected, sky subtracted and mosaicked using
the ORAC-DR pipeline.1 WFCAM data were processed using the
WFCAM science pipeline by the Cambridge Astronomical Surveys
Unit (CASU) (Irwin et al. 2004), and archived at the WSA (Hambly
et al. 2008). Observations consisted of a three-point jitter pattern in
the Y and J bands, and five-point jitter patterns in the H and K bands
repeated twice. All data were acquired with 2 × 2 microstepping.
The WFCAM and UFTI filters are on the Mauna Kea Observatories
(MKO) photometric system (Tokunaga, Simons & Vacca 2002).
The majority of our z photometry was taken using the Device Optimized for the LOw RESolution (DOLORES; Molinari, Conconi
& Pucillo 1997) at the Telescopio Nazionale Galileo (TNG). The
observations were taken under programme AOT22 TAC 96 spanning from 2010 to 2012. DOLORES is equipped with a 2048 ×
2048 pixels CCD with a field of view of 8.6 × 8.6 arcmin with a
0.252 arcsec pixel−1 scale. The observations were taken with the
z Sloan filter. A small number of targets were observed in the
z band using the using the ESO Faint Object Spectrograph and
Camera (EFOSC2) mounted on the New Technology Telescope
(NTT; programme 082.C-0399) and using the Auxiliary-port Camera (ACAM) on the William Herschel Telescope (WHT). For each
epoch a set of standard calibration flat-fields and darks observations
were taken. The images were dark subtracted, flat-fielded and in the
case of multiple exposures combined using standard IRAF routines.
The data were taken in different observing conditions, from photometric conditions to cirrus. No attempts to perform defringing
to the images were made. The E2V4240 CCD detector in use in
DOLORES has a low fringing level, the science object was also
normally located in the top-right section of the CCD where fringing
is even smaller. Photometry was performed with IRAF using a fixed
circular aperture with radius 2 arcsec. The photometric zero-point
was calibrated using the non-saturated SDSS stars present in the
field of view.
The SDSS z band filter is slightly peculiar in that it has no red
cut-off. Instead, the red cut-off is defined by the detector sensitivity. So, although the DOLORES and ACAM data were taken
through an SDSS z band filter, this does not trivially lead to the
photometry being on the SDSS system, since there may be differences in the detectors’ long-wavelength responses. To check
the consistency of the DOLORES and ACAM photometric systems with SDSS, we have compared synthetic photometry for a
set of reference stellar spectra convolved with the combined filter
and detector response curves for each of the systems. The difference was found to be much smaller than the typical scatter in the
zero-point from the reference stars, which dominate our quoted
errors, and so we did not correct the SDSS reference stars magnitudes before calibration. For objects as red as T dwarfs, however,
the difference can be more significant. Synthesized photometry using template-T dwarf spectra found a mean offset of close to zero
for DOLORES (z (DOLORES) − z (SDSS) = −0.02 ± 0.02). For
ACAM, we found a small offset of z (ACAM) − z (SDSS) =
+0.09 ± 0.03.
For the EFOSC2 observations, a Gunn z-band filter (ESO Z#623)
was used and we used the transform given in Burningham et al.
(2009) to calculate zEFOSC2 for the SDSS secondary calibrators. To
place the resulting zEFOSC2 (AB) photometry on the Sloan z (AB)
system, we used the transform determined in Burningham et al.
(2010b): z(EFOSC2) − z (SDSS) = −0.19 ± 0.02.
The best available broad-band photometry for all targets presented here is given in Table 1.
α
(J2000)
00:07:34.90
01:27:35.66
01:28:55.07
01:30:17.79
01:33:02.48
01:39:50.51
02:00:13.18
02:26:03.18
02:45:57.88
02:55:45.28
03:29:20.22
07:45:02.79
07:46:16.98
07:47:20.07
07:58:29.83
07:59:37.75
08:00:48.27
08:09:18.41
08:11:10.86
08:14:07.51
08:15:07.26
08:19:18.58
08:21:55.49
08:47:43.93
09:26:08.82
09:27:44.20
09:29:06.75
09:32:45.48
09:50:47.28
09:54:29.90
10:21:44.87
10:23:05.44
10:29:40.52
10:42:24.20
10:43:55.37
10:51:34.32
10:53:34.64
11:11:27.77
11:37:17.17
11:52:29.68
11:52:39.94
11:55:08.39
12:04:44.67
ULAS J000734.90+011247.1
ULAS J012735.66+153905.9
ULAS J012855.07+063357.0
ULAS J013017.79+080453.9
CFBDS J013302+023128a
ULAS J013950.51+150307.6
ULAS J020013.18+090835.2
ULAS J022603.18+070231.4
ULAS J024557.88+065359.4
ULAS J025545.28+061655.8
ULAS J032920.22+043024.5
ULAS J074502.79+233240.3
ULAS J074616.98+235532.2
ULAS J074720.07+245516.3
ULAS J075829.83+222526.7
ULAS J075937.75+185555.0
ULAS J080048.27+190823.8
ULAS J080918.41+212615.2
ULAS J081110.86+252931.8
ULAS J081407.51+245200.9
ULAS J081507.26+271119.2
ULAS J081918.58+210310.4
ULAS J082155.49+250939.6
ULAS J084743.93+035040.2
ULAS J092608.82+040239.7
ULAS J092744.20+341308.7
WISEP J092906.77+040957.9b
ULAS J093245.48+310206.4
ULAS J095047.28+011734.3
ULAS J095429.90+062309.6
ULAS J102144.87+054446.1
ULAS J102305.44+044739.2
ULAS J102940.52+093514.6
ULAS J104224.20+121206.8
ULAS J104355.37+104803.4
ULAS J105134.32−015449.8
ULAS J105334.64+015719.7
ULAS J111127.77+051855.5
ULAS J113717.17+112657.2
ULAS J115229.68+035927.3
ULAS J115239.94+113407.6
ULAS J115508.39+044502.3
ULAS J120444.67−015034.9
z
–
–
22.73 ± 0.40A
–
–
–
–
–
–
–
20.75 ± 0.17
–
>20.99D
–
–
23.32 ± 0.09D
22.05 ± 0.16D
–
–
>21.05D
–
21.93 ± 0.08E
–
21.90 ± 0.10A
–
>21.8D
–
–
–
–
–
23.50 ± 0.17D
–
–
–
–
–
–
>22.24D
–
–
–
–
δ
(J2000)
+01:12:47.10
+15:39:05.90
+06:33:57.00
+08:04:53.90
+02:31:28.90
+15:03:07.60
+09:08:35.20
+07:02:31.40
+06:53:59.40
+06:16:55.80
+04:30:24.50
+23:32:40.30
+23:55:32.20
+24:55:16.30
+22:25:26.70
+18:55:55.00
+19:08:23.80
+21:26:15.20
+25:29:31.80
+24:52:00.90
+27:11:19.20
+21:03:10.40
+25:09:39.60
+03:50:40.20
+04:02:39.70
+34:13:08.70
+04:09:57.70
+31:02:06.40
+01:17:34.30
+06:23:09.60
+05:44:46.10
+04:47:39.20
+09:35:14.60
+12:12:06.80
+10:48:03.40
−01:54:49.80
+01:57:19.70
+05:18:55.50
+11:26:57.20
+03:59:27.30
+11:34:07.60
+04:45:02.30
−01:50:34.90
19.22 ± 0.07
19.47 ± 0.13
19.66 ± 0.14
19.06 ± 0.03W
19.36 ± 0.11
19.72 ± 0.17
18.98 ± 0.07
19.62 ± 0.05W
19.43 ± 0.1
19.15 ± 0.07
18.55 ± 0.02W
20.0 ± 0.15
20.18 ± 0.19
19.35 ± 0.05W
18.68 ± 0.04
20.21 ± 0.18
19.76 ± 0.12
19.65 ± 0.09
18.76 ± 0.03
19.6 ± 0.1
19.48 ± 0.1
18.25 ± 0.03
18.61 ± 0.04
19.61 ± 0.05W
19.7 ± 0.09
19.66 ± 0.14
17.89 ± 0.01W
20.0 ± 0.09
18.9 ± 0.03W
17.73 ± 0.01W
18.82 ± 0.03W
19.49 ± 0.05W
18.24 ± 0.02W
19.58 ± 0.09
19.21 ± 0.03W
18.85 ± 0.03W
19.77 ± 0.1
19.87 ± 0.1
20.14 ± 0.21
18.54 ± 0.03
19.3 ± 0.06
19.38 ± 0.07
17.99 ± 0.03
YMKO
18.05 ± 0.04
18.22 ± 0.07
18.93 ± 0.12
17.93 ± 0.02W
18.34 ± 0.08
18.44 ± 0.1
17.81 ± 0.04
18.52 ± 0.04W
18.36 ± 0.04W
18.04 ± 0.03W
17.55 ± 0.02W
18.88 ± 0.07
19.0 ± 0.08
18.17 ± 0.05W
17.62 ± 0.02W
18.7 ± 0.07
18.55 ± 0.06
18.58 ± 0.03W
17.57 ± 0.02
18.54 ± 0.05
18.31 ± 0.03W
16.95 ± 0.01
17.23 ± 0.01W
18.53 ± 0.04W
18.59 ± 0.06
18.77 ± 0.11
16.87 ± 0.01W
18.73 ± 0.05
18.02 ± 0.03W
16.6 ± 0.01W
17.66 ± 0.02W
18.39 ± 0.04W
17.28 ± 0.01W
18.52 ± 0.06
18.23 ± 0.02W
17.75 ± 0.02W
18.5 ± 0.06
18.74 ± 0.07
18.5 ± 0.09
17.28 ± 0.02
18.26 ± 0.04
18.33 ± 0.05
16.74 ± 0.02U
JMKO
–
18.62 ± 0.14
–
18.21 ± 0.02W
18.51 ± 0.15
18.53 ± 0.18
18.18 ± 0.11
18.82 ± 0.03W
18.95 ± 0.17
18.4 ± 0.02W
17.89 ± 0.02W
–
–
18.5 ± 0.04W
17.91 ± 0.02W
–
–
18.99 ± 0.03W
18.19 ± 0.12
–
18.6 ± 0.03W
17.28 ± 0.03
17.24 ± 0.01W
18.71 ± 0.03W
–
–
17.24 ± 0.01W
19.04 ± 0.23
18.4 ± 0.03W
16.87 ± 0.01W
17.96 ± 0.02W
18.73 ± 0.04W
17.63 ± 0.01W
18.9 ± 0.12
18.58 ± 0.02W
18.07 ± 0.02W
–
–
–
17.7 ± 0.05
18.66 ± 0.1
–
17.1 ± 0.02U
HMKO
–
–
–
18.35 ± 0.04W
–
–
18.18 ± 0.2
18.79 ± 0.06W
–
–
18.4 ± 0.04W
–
–
18.53 ± 0.07W
17.87 ± 0.12
–
–
18.65 ± 0.22
18.02 ± 0.19
–
–
17.18 ± 0.06
17.23 ± 0.09
18.99 ± 0.08W
–
–
17.61 ± 0.02W
–
18.85 ± 0.07W
17.05 ± 0.01W
17.97 ± 0.03W
18.58 ± 0.07W
17.64 ± 0.02W
–
18.66 ± 0.05W
18.27 ± 0.04W
–
–
–
17.77 ± 0.12
18.32 ± 0.17
–
17.29 ± 0.09
KMKO
−0.91 ± 0.13
−0.88 ± 0.17
−0.81 ± 0.14
−0.72 ± 0.07
<−1.19
−0.83 ± 0.11
−0.83 ± 0.11
–
–
–
–
−1.62 ± 0.17
−0.87 ± 0.17L
–
–
−0.95 ± 0.12
−0.73 ± 0.28L
–
−1.03 ± 0.13
−0.30 ± 0.10L
–
–
–
−0.65 ± 0.14
−0.69 ± 0.14
−1.27 ± 0.28
−0.92 ± 0.07
–
–
−0.58 ± 0.09
−0.99 ± 0.26
−0.52 ± 0.12L
−1.56 ± 0.17
−0.93 ± 0.15
−1.36 ± 0.22
−0.56 ± 0.14
−1.10 ± 0.16
–
−0.51 ± 0.16
−0.69 ± 0.06
<−1.38
–
–
CH4 s − CH4 l
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Name
T6.1
T6.0
T5.8
T5.5
N/A
T5.9
T5.9
N/A
N/A
N/A
N/A
T7.6
N/A
N/A
N/A
T6.2
N/A
N/A
T6.4
N/A
N/A
N/A
N/A
T5.3
T5.4
T7.0
T6.1
N/A
N/A
T5.0
T6.3
N/A
T7.5
T6.2
T7.2
T4.9
T6.6
N/A
T4.7
T5.4
N/A
N/A
N/A
T5.8
T5.5
T5.4
T5.3
N/A
T5.6
T5.5
N/A
N/A
N/A
N/A
T7.4
N/A
N/A
N/A
T5.9
N/A
N/A
T6.1
N/A
N/A
N/A
N/A
T4.7
T4.9
T6.4
T6.0
N/A
N/A
T4.6
T5.6
N/A
T7.3
T5.8
T6.7
T4.4
T6.2
N/A
T4.0
T5.2
N/A
N/A
N/A
T6.4
T6.4
T6.2
T5.7
N/A
T6.2
T6.1
N/A
N/A
N/A
N/A
T7.9
N/A
N/A
N/A
T6.5
N/A
N/A
T6.7
N/A
N/A
N/A
N/A
T5.7
T5.8
T7.5
T6.3
N/A
N/A
T5.3
T6.9
N/A
T7.8
T6.5
T7.5
T5.4
T6.9
N/A
T5.3
T5.6
N/A
N/A
N/A
T7
T6.5
T6
T6
T8
T7
T6
T7
T7
T6
T5
T9d
T7
T6.5
T6.5
T6
N/A
T8
T7
T5p
T7p
T6
T4.5
N/A
T6
T5.5
T7
T2
T8p
T5
T6
T6.5
T8
T7.5p
T8
T6
T6.5
T4.5
N/A
T6
T8.5d
T7
T4.5
CH4 type CH4 type CH4 type Spectral
(min)
(max)
type
Table 1. Best available NIR photometry for our sample. No superscript on a broad-band photometric value indicates UKIDSS survey photometry for YJHK, SDSS DR8 for z band. Unless indicated otherwise, all
CH4 photometry is from TNG/NICS. Superscripts refer to the following instruments: A = ACAM (WHT); D = DOLORES (TNG); E = EFOSC2 (NTT); L = LIRIS(WHT); U = UFTI (UKIRT); W = WFCAM
(UKIRT). z band photometry has been converted to the SDSS system as described in the text.
76 T dwarfs
461
12:06:21.03
12:12:26.80
12:23:43.35
12:54:46.35
12:58:35.97
12:59:39.44
13:02:27.54
13:35:02.11
13:38:28.69
13:39:33.64
13:39:43.79
14:17:56.22
14:21:45.63
14:25:36.35
14:49:02.02
15:16:37.89
15:17:21.12
15:34:06.06
15:36:53.80
15:49:14.45
16:01:43.75
16:14:36.96
16:17:10.39
16:19:34.78
16:19:38.12
16:26:55.04
16:39:31.52
21:16:16.26
22:37:28.91
23:00:49.08
23:15:36.93
23:18:56.24
23:26:00.40
23:26:24.07
23:31:04.12
23:42:28.97
23:52:04.62
23:57:15.98
ULAS J120621.03+101802.9
ULAS J121226.80+101007.4
ULAS J122343.35-013100.7
ULAS J125446.35+122215.7
ULAS J125835.97+030736.1
ULAS J125939.44+293322.4
ULAS J130227.54+143428.0
ULAS J133502.11+150653.5
ULAS J133828.69−014245.4
ULAS J133933.64−005621.1
ULAS J133943.79+010436.4
ULAS J141756.22+133045.8
ULAS J142145.63+013619.0
ULAS J142536.35+045132.3
ULAS J144902.02+114711.4
ULAS J151637.89+011050.1
WISE J151721.13+052929.3c
ULAS J153406.06+055643.9
ULAS J153653.80+015540.6
ULAS J154914.45+262145.6
ULAS J160143.75+264623.4
ULAS J161436.96+244230.1
ULAS J161710.39+235031.4
ULAS J161934.78+235829.3
ULAS J161938.12+300756.4
ULAS J162655.04+252446.8
ULAS J163931.52+323212.7
ULAS J211616.26-010124.3
ULAS J223728.91+064220.1
ULAS J230049.08+070338.0
ULAS J231536.93+034422.7
ULAS J231856.24+043328.5
ULAS J232600.40+020139.2
ULAS J232624.07+050931.6
ULAS J233104.12+042652.6
ULAS J234228.97+085620.1
ULAS J235204.62+124444.9
ULAS J235715.98 +0:3:40.
z
–
–
–
–
–
–
>19.12D
–
–
–
–
20.42 ± 0.16
–
>21.87D
–
–
–
22.52 ± 0.15D
–
–
21.35 ± 0.05D
22.36 ± 0.35D
–
–
–
–
20.30 ± 0.11
>22.10D
–
21.6 ± 0.12D
–
25.6 ± 0.32D
–
21.85 ± 0.21D
22.12 ± 0.26D
20.15 ± 0.12
–
21.42 ± 0.22D
δ
(J2000)
+10:18:02.90
+10:10:07.40
−01:31:00.70
+12:22:15.70
+03:07:36.10
+29:33:22.40
+14:34:28.00
+15:06:53.50
−01:42:45.40
−00:56:21.10
+01:04:36.40
+13:30:45.80
+01:36:19.00
+04:51:32.30
+11:47:11.40
+01:10:50.10
+05:29:29.03
+05:56:43.90
+01:55:40.60
+26:21:45.60
+26:46:23.40
+24:42:30.10
+23:50:31.40
+23:58:29.30
+30:07:56.40
+25:24:46.80
+32:32:12.70
−01:01:24.30
+06:42:20.10
+07:03:38.00
+03:44:22.70
+04:33:28.50
+02:01:39.20
+05:09:31.60
+04:26:52.60
+08:56:20.10
+12:44:44.90
+01:32:40.30
20.57 ± 0.23
20.48 ± 0.25
19.71 ± 0.13
19.51 ± 0.11
19.7 ± 0.14
19.65 ± 0.09
19.75 ± 0.13
19.03 ± 0.03U
19.57 ± 0.08W
19.21 ± 0.05W
19.15 ± 0.05W
17.94 ± 0.03
19.31 ± 0.12
20.02 ± 0.14
18.35 ± 0.04
19.48 ± 0.12
19.57 ± 0.07
20.24 ± 0.19
19.15 ± 0.08
19.15 ± 0.07
19.48 ± 0.08W
19.42 ± 0.08
18.99 ± 0.05
19.72 ± 0.11
19.84 ± 0.11
19.82 ± 0.11
18.14 ± 0.02
19.53 ± 0.12
19.79 ± 0.08W
18.97 ± 0.04W
19.89 ± 0.12
20.18 ± 0.13
19.4 ± 0.08
19.75 ± 0.15
20.16 ± 0.14
17.37 ± 0.01W
19.64 ± 0.11
19.78 ± 0.06W
b Kirkpatrick
YMKO
19.11 ± 0.15W
18.69 ± 0.09W
18.7 ± 0.09
18.29 ± 0.06
18.38 ± 0.05W
18.39 ± 0.06
18.6 ± 0.04W
17.97 ± 0.02U
18.69 ± 0.1W
18.24 ± 0.05W
18.08 ± 0.04
16.77 ± 0.01
18.52 ± 0.04W
18.7 ± 0.09
17.36 ± 0.02
18.41 ± 0.05W
18.54 ± 0.05
19.02 ± 0.1
17.93 ± 0.05
18.05 ± 0.03W
18.43 ± 0.05W
18.52 ± 0.04
17.72 ± 0.02
18.62 ± 0.06W
18.61 ± 0.07W
18.4 ± 0.04W
16.71 ± 0.01
18.27 ± 0.07
18.78 ± 0.05W
17.67 ± 0.02W
18.79 ± 0.08
18.78 ± 0.07
17.98 ± 0.04
18.61 ± 0.1
18.67 ± 0.08
16.39 ± 0.01W
18.27 ± 0.05
18.5 ± 0.04W
JMKO
19.53 ± 0.09W
19.06 ± 0.08W
–
18.62 ± 0.17
18.59 ± 0.05W
18.55 ± 0.14
18.8 ± 0.04W
18.3 ± 0.03U
19.14 ± 0.09W
18.48 ± 0.04W
18.39 ± 0.13
17.0 ± 0.03
18.54 ± 0.03W
–
17.73 ± 0.07
18.67 ± 0.06W
18.85 ± 0.15
–
18.03 ± 0.1
18.29 ± 0.03W
18.82 ± 0.07W
–
18.16 ± 0.08
18.91 ± 0.06W
18.79 ± 0.06W
18.62 ± 0.04W
16.72 ± 0.03
–
19.23 ± 0.04W
17.77 ± 0.03W
–
–
18.46 ± 0.12
18.61 ± 0.14
–
16.77 ± 0.01W
18.55 ± 0.16
18.68 ± 0.03W
HMKO
–
–
–
18.26 ± 0.2
–
–
–
18.23 ± 0.14
19.21 ± 0.12W
18.39 ± 0.05W
18.39 ± 0.05W
17.0 ± 0.04
–
–
18.1 ± 0.15
18.49 ± 0.2
–
–
18.01 ± 0.16
–
18.75 ± 0.08W
–
–
–
–
–
16.8 ± 0.06
–
19.94 ± 0.18W
17.74 ± 0.05W
–
–
18.41 ± 0.2
–
–
17.1 ± 0.02W
18.41 ± 0.21
18.6 ± 0.05W
KMKO
–
–
−0.66 ± 0.16
−0.56 ± 0.19
–
−0.59 ± 0.13
−0.23 ± 0.10
−0.60 ± 0.09
–
–
–
−0.51 ± 0.08
–
−0.93 ± 0.12
−0.48 ± 0.13
−0.96 ± 0.20
–
−0.56 ± 0.13
–
−0.60 ± 0.12
–
−1.01 ± 0.15
−0.70 ± 0.09
−0.75 ± 0.11
−0.43 ± 0.09
–
–
−1.10 ± 0.31
–
−0.43 ± 0.08
< -0.77
< -0.95
−1.64 ± 0.16
−0.51 ± 0.12
–
–
−0.95 ± 0.12
–
N/A
N/A
T5.3
T4.9
N/A
T5.0
T3.1
T5.1
N/A
N/A
N/A
T4.7
N/A
T6.2
T4.6
T6.3
N/A
T4.9
N/A
T5.1
N/A
T6.4
T5.4
T5.6
T4.3
N/A
N/A
T6.6
N/A
T4.3
N/A
N/A
T7.7
T4.7
N/A
N/A
T6.2
N/A
N/A
N/A
T4.7
T4.1
N/A
T4.5
T2.4
T4.8
N/A
N/A
N/A
T4.4
N/A
T5.9
T4.0
T5.7
N/A
T4.4
N/A
T4.6
N/A
T6.0
T5.1
T5.3
T3.9
N/A
N/A
T5.8
N/A
T3.9
N/A
N/A
T7.5
T4.2
N/A
N/A
T5.9
N/A
N/A
N/A
T5.8
T5.6
N/A
T5.4
T3.7
T5.4
N/A
N/A
N/A
T5.0
N/A
T6.4
T5.1
T6.7
N/A
T5.3
N/A
T5.5
N/A
T6.7
T5.7
T5.9
T4.7
N/A
N/A
T7.2
N/A
T4.7
N/A
N/A
T7.9
T5.1
N/A
N/A
T6.5
N/A
T5
T5
T6
N/A
T5
T5
T4.5
T6
T7.5
T7
T5
T5
T4.5
T6.5
T5.5
T6.5
T8p
T5
T5
T5
T6.5
T7
T6
T6
T5
T5
T3
T6
T6.5p
T4.5
T7
T7.5
T8
N/A
T4
T6
T6.5
T5.5p
CH4 s − CH4 l CH4 type CH4 type CH4 type Spectral
(min)
(max)
type
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
et al. (2011); UKIDSS designation: ULAS J013302.48+023128.9.
et al. (2011); UKIDSS designation: ULAS J092906.75+04:0957.7.
c Mace et al. (2013); UKIDSS designation: ULAS J151721.12+052929.0.
d On the spectral typing system of Burningham et al. (2008).
a Albert
α
(J2000)
Name
Table 1 – continued
462
B. Burningham et al.
76 T dwarfs
4 S P E C T RO S C O P I C C O N F I R M AT I O N
Spectroscopic confirmation of most T dwarf candidates that survived the photometric follow-up programme was achieved using
the Near InfraRed Imager and Spectrometer (NIRI; Hodapp et al.
2003) and the Gemini Near Infrared Spectrograph (GNIRS; Elias
et al. 2006) on the Gemini North Telescope5 and the InfraRed Camera and Spectrograph (IRCS; Kobayashi et al. 2000) on the Subaru
telescope, both on Mauna Kea, Hawaii. In addition a smaller number
of spectra were obtained using the Folded port InfraRed Echellette
(FIRE) spectrograph (Simcoe et al. 2008, 2010) mounted on the
Baade 6.5 m Magellan telescope at Las Campanas Observatory. We
also obtained spectroscopy for a single target using XSHOOTER
(Vernet et al. 2011) on UT2 of the VLT (Programme ID: 086.C0450).
All observations were made up of a set of subexposures in an
ABBA jitter pattern to facilitate effective background subtraction,
with a slit width of 1 arcsec for NIRI, GNIRS and IRCS, whilst
0.6 arcsec was used for the FIRE observations. The length of the
A–B jitter was 10 arcsec. For targets brighter than J = 18.5 total
integrations were typically 4 × 300 s for NIRI, GNIRS and IRCS
observations, whilst fainter targets were typically integrated for 8 ×
300 s. FIRE integrations were 2 × 120 s for J < 18.0, 4 × 150 s
for 18.0 < J < 18.5, 6 × 150 s for 18.50 < J < 18.6 and 8 ×
150 s for our faintest targets. The programme numbers and dates of
individual observations are summarized in Appendix C.
The NIRI and GNIRS observations were reduced using standard IRAF Gemini packages (Cooke & Rodgers 2005). The Subaru
IRCS spectra were extracted using standard IRAF packages. The AB
pairs were subtracted using generic IRAF tools, and median stacked.
The NIRI, GNIRS and IRCS spectra were calibrated in a similar
4
http://www.iac.es/galeria/jap/lirisdr/LIRIS_DATA_REDUCTION.html
under programmes GN-2009A-Q-16, GN-2009B-Q-62, GN-2009B-Q99, GN-2010A-Q-44, GN-2010B-Q-41, GN-2011A-Q-73, GN-2011B-Q-5,
GN-2011B-Q-43 and GN-2012A-Q-84
5
manner. Comparison argon arc frames were used to obtain dispersion solutions, which were then applied to the pixel coordinates in
the dispersion direction on the images. The resulting wavelengthcalibrated subtracted pairs had a low-level of residual sky emission removed by fitting and subtracting this emission with a set
of polynomial functions fitted to each pixel row perpendicular to
the dispersion direction, and considering pixel data on either side
of the target spectrum only. The spectra were then extracted using
a linear aperture, and cosmic rays and bad pixels removed using
a sigma-clipping algorithm. Telluric correction was achieved by
dividing each extracted target spectrum by that of an early A- or
F-type standard star observed just before or after the target and at
a similar airmass. Prior to division, hydrogen lines were removed
from the standard star spectrum by interpolating the stellar continuum. Relative flux calibration was then achieved by multiplying
through by a blackbody spectrum of the appropriate Teff . This Teff
was taken from Masana, Jordi & Ribas (2006) where available, or
else was estimated from the spectral type of the telluric standard.
Since the near-infrared region is well into the Rayleigh–Jeans tail
of an A or F star’s spectrum, very little systematic error is likely to
be introduced from a crude estimate of this Teff .
The FIRE spectra were extracted using the low-dispersion version
of the FIREHOSE pipeline, which is based on the MASE pipeline
(Bochanski et al. 2009; Simcoe et al. 2010). The pipeline uses a
flat-field constructed from two quartz lamp images taken with the
lamp at high (2.5 V) and low (1.5 V) voltage settings. The data were
divided by this pixel flat before being wavelength calibrated. The
pipeline performs sky subtraction following the method outlined
in Bochanski et al. (2011), adapted for the low-dispersion configuration of the spectrograph. The spectra were optimally extracted
before being combined using a weighted mean, using an adaptation of the xcombspec routine from SPEXTOOL (Cushing, Vacca &
Rayner 2004). The T dwarf spectra were then corrected for telluric
absorption and flux calibrated using a FIRE-specific version of the
xtellcor routine (Vacca, Cushing & Rayner 2003). Finally, residual
outlying points due to cosmic rays and bad pixels were removed
using a simple 3σ clipping algorithm.
The XSHOOTER data were reduced using the ESO pipeline
(version 1.3.7). The pipeline removes non-linear pixels, subtracts
the bias (in the VIS arm) or dark frames (in the NIR arm) and divides
the raw frames by flat-fields. Images are pair-wise subtracted to
remove sky background. The pipeline then extracts and merges the
different orders in each arm, rectifying them using a multipinhole
arc lamp (taken during the day-time calibration) and correcting for
the flexure of the instrument using single-pinhole arc lamps (taken
at night, one for each object observed). Telluric stars are reduced
in the same way, except that sky subtraction is done by fitting
the background (as tellurics are not observed in nodding mode).
The spectra were telluric corrected and flux calibrated using IDL
routines, following a standard procedure: first the telluric spectrum
is cleared of H I absorption lines (by interpolating over them) and
scaled to match the measured magnitudes; then is divided by a
blackbody curve for the appropriate temperature (estimated from
the telluric standard’s spectral type), to obtain the instrument +
atmosphere response curve; finally the target spectra are multiplied
by the response curve obtained to flux calibrate it. The arms (VIS and
NIR) were then merged by matching the flux level in the overlapping
regions between them.
Complete details of the spectroscopic observations obtained for
each of the T dwarfs presented here are given in Appendix C. The
resulting spectra are shown in Fig. 3. This includes one T dwarf
(ULAS J0929+0409) confirmed by Kirkpatrick et al. (2011), three
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
An estimate of the spectral type was obtained using the conversion defined by equation 2 from Tinney et al. (2005). The resulting methane colours for all targets with spectra presented here are
given in Table 1, along with spectroscopically determined spectral types (see Section 4) and photometric spectral types. We also
present CH4 photometry for several DR8 targets which have CH4 s −
CH4 l < −0.5, but which for various reasons were not followed up
with spectroscopy. These targets are included as they form part of
our UKIDSS DR8 space density estimate outlined in Section 9. A
full summary of all the CH4 photometry obtained, including earlier
type objects and the extension of our analysis to include photometrically confirmed T dwarfs in DR9, along with a more detailed
description of the CH4 calibration can be found in Cardoso et al. (in
preparation).
For a small number of targets, we obtained differential methane
photometry using the Long-slit Infrared Imaging Spectrograph
(LIRIS; Manchado et al. 1998) mounted on the WHT. These
data were flat-field corrected, sky subtracted and mosaicked using
LIRIS-DR.4 In these cases, the methane colours are constructed as
H − [CH4 ]l, and calibrated assuming that the average H − [CH4 ]l of
bright secondary calibrators in the field was zero (see also Kendall
et al. 2007; Pinfield et al. 2008). Since no calibration for spectral
type is yet determined for the H − CH4 l colour, we do not present
photometric estimates for the spectral types from these data.
463
464
B. Burningham et al.
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Figure 3. Spectra of the 76 T dwarfs presented here. Each spectrum is normalized at 1.27 ± 0.005 µm and offset for clarity. The GNIRS, NIRI and IRCS
spectra have been rebinned by a factor of 3 as a compromise to maximize S/N and not sacrifice resolution.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 3 – continued
465
76 T dwarfs
B. Burningham et al.
466
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Figure 3 – continued
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 3 – continued
467
76 T dwarfs
B. Burningham et al.
468
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Figure 3 – continued
76 T dwarfs
469
(ULAS J0954+0623, ULAS J1204-0150, ULAS J1152+0359)
confirmed by Scholz et al. (2012) and one (ULAS J1517+0529)
confirmed by Mace et al. (2013) since our spectroscopic follow-up,
but prior to this publication.
for the T spectral sequence (Burningham et al. 2010b; Leggett et al.
2010a; Liu et al. 2012).
5 S PAC E - BA S E D M I D - I R P H OT O M E T RY
4.1 Spectral types
We have assigned spectral types following the scheme of Burgasser
et al. (2006) for types as late as T8, and the extension of Burningham et al. (2008) for types beyond T8. We have adopted this scheme
in this work for two reasons. First, it provides continuity with our
previous work (Burningham et al. 2008, 2009, 2010b), allowing
a meaningful update to our previous space density estimate. Secondly, as was discussed in detail in Burningham et al. (2008), this
scheme provides excellent continuity with the evolution of spectral
index values from earlier types. This does not diminish the scheme’s
fundamentally empirical nature; it is anchored to template objects.
However, it does seek to minimize the subjectivity as to the degree of
spectral difference required for the distinction of two subtypes. We
have only identified two new T8+ dwarfs in this paper, so this issue
is of minor importance, and we have indicated the spectral types of
these two objects on the Cushing et al. (2011) extension scheme in
the notes column for completeness, and to avoid any future confusion. Objects that show substantial discrepancy either between the
spectral types indicated by their spectral indices, or in comparison
to their best-fitting spectral template have been classified as peculiar, and are denoted with the suffix ‘p’. The most common feature
leading to this designation is the suppression or enhancement of
the K-band peak relative to the template. Table 2 summarizes our
spectral type measurements and adopted classifications.
In Fig. 4, we have plotted our full UKIDSS LAS sample of 146
spectroscopically confirmed T dwarfs with YJH photometry (from
a total of 171) on a YJH colour–colour diagram. Spectral types are
distinguished by coloured symbols in whole subtype bins. Although
there is significant scatter, a general trend from top right to lower
left is apparent for the T4–T8 dwarfs, with the earliest type objects
dominating the top right of the plot, whilst the later type objects
dominate the lower left. This is consistent with the reduced H-band
flux due to deepening CH4 absorption, and a general trend to bluer
Y − J colours with decreasing Teff that have previously been noted
5.1 WISE cross-matched photometry
We cross-matched our full list of 171 spectroscopically confirmed
T dwarfs within the UKIDSS LAS (which includes some objects
that were confirmed in the literature rather than by our follow-up)
against the WISE all-sky release catalogue, with matching radius of
6 arcsec. The typical epoch difference between the UKIDSS and
WISE observations is less than 3 yr, so this ensured that high proper
motion targets would still be matched. All apparent matches were
visually inspected to remove spurious correlations. 67 T dwarfs
were found with WISE photometry, of which 6 were affected by
blending with another source. 25 of the 67 T dwarfs with WISE
photometry are confirmed here for the first time, and their details
are given in Table 3.
In Figs 5 and 6, we compare the numbers of T dwarfs detected by
UKIDSS with those detected by both WISE and UKIDSS for objects
brighter than J = 18.3 and for objects in the J = 18.3–18.8 range
(the faintest 0.5 mag bin of our complete sample). It can be seen
that UKIDSS is considerably more sensitive to earlier type objects
than WISE, with roughly twice as many T4–T6.5 dwarfs identified
in the J < 18.3 regime. In the fainter 18.3 < J < 18.8 regime,
this effect is even more pronounced, and extends to the T7–T7.5
bin. This highlights that although WISE now dominates the search
for cool and faint T8+ dwarfs, wide and deep NIR surveys such
as UKIDSS, and the (wider) VHS and (deeper) VIKING surveys
continue represent an important resource for exploring the L and T
dwarf sequences.
The faint nature of the large number of L and T dwarfs that will
be revealed by VHS and VIKING, in particular, will be extremely
challenging to confirm spectroscopically. To take advantage of the
opportunity these offer for robust statistical studies of the substellar
component of the Galaxy, it will be essential to develop methods
for determining their properties from the photometric data that will
be supplied by the surveys.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 3 – continued
Adopted
T7
T6.5
T6
T6
T8
T7
T6
T7
T7
T6
T5
T9
T7
T6.5
T6.5
T6
T8
T7
T5p
T7p
T6
T4.5
T6
T5.5
T7
T2
T8p
T5
T6
T6.5
T8
T7.5p
T8
T6
T6.5
T4.5
T6
T8.5
T7
T4.5
T5
T5
ULAS J0007+0112
ULAS J0127+1539
ULAS J0128+0633
ULAS J0130+0804
CFBDS J0133+0231
ULAS J0139+1503
ULAS J0200+0908
ULAS J0226+0702
ULAS J0245+0653
ULAS J0255+0616
ULAS J0329+0430
ULAS J0745+2332
ULAS J0746+2355
ULAS J0747+2455
ULAS J0758+2225
ULAS J0759+1855
ULAS J0809+2126
ULAS J0811+2529
ULAS J0814+2452
ULAS J0815+2711
ULAS J0819+2103
ULAS J0821+2509
ULAS J0926+0402
ULAS J0927+3413
WISEP J0929+0409
ULAS J0932+3102
ULAS J0950+0117
ULAS J0954+0623
ULAS J1021+0544
ULAS J1023+0447
ULAS J1029+0935
ULAS J1042+1212
ULAS J1043+1048
ULAS J1051−0154
ULAS J1053+0157
ULAS J1111+0518
ULAS J1152+0359
ULAS J1152+1134
ULAS J1155+0445
ULAS J1204−0150
ULAS J1206+1018
ULAS J1212+1010
T7
T6.5
T6
T6
T8
T7
T6
T7
T7
T6
T5
T9
T7
T6.5
T6.5
T6
T8
T7
T5
T7
T6
T4.5
T6
T5.5
T7
T2
T8p
T5
T6
T6.5
T8
T7.5p
T8
T6
T6.5
T4.5
T6
T8.5
T7
T4.5
T5
T5
Templ.
0.103
0.179
0.142
0.170
0.051
0.122
0.146
0.082
0.101
0.154
0.238
0.024
0.111
0.126
0.130
0.191
0.026
0.113
0.237
0.112
0.183
0.326
0.186
0.200
0.087
0.431
0.043
0.200
0.167
0.100
0.049
0.064
0.047
0.146
0.092
0.292
0.142
0.038
0.126
0.348
0.227
0.333
0.007 (T6/7)
0.007 (T6)
0.008 (T6)
0.005 (T6)
0.018 (>T7)
0.014 (T6/7)
0.008 (T6/7)
0.004 (T7)
0.003 (T7)
0.002 (T6/7)
0.007 (T4/5)
0.039 (>T7)
0.009 (T7)
0.005 (T6/7)
0.006 (T6)
0.004 (T6)
0.003 (>T7)
0.002 (T7)
0.007 (T5)
0.002 (T7)
0.004 (T6)
0.004 (T4)
0.007 (T6)
0.011 (T6)
0.002 (T7)
0.027 (T2)
0.003 (T7)
0.003 (T5)
0.002 (T6)
0.010 (T6/7)
0.001 (>T7)
0.007 (T6/7)
0.003 (T7)
0.002 (T6)
0.029 (T7)
0.051 (T5/6)
0.001 (T6)
0.004 (>T7)
0.017 (T6)
0.008 (T5)
0.020 (T5)
0.022 (T4/5)
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.286
0.307
0.320
0.338
0.191
0.268
0.285
0.230
0.260
0.282
0.452
0.130
0.226
0.284
0.340
0.324
0.193
0.259
0.408
0.243
0.350
0.508
0.317
0.347
0.276
0.643
0.218
0.370
0.331
0.279
0.182
0.286
0.221
0.354
0.245
0.408
0.345
0.147
0.319
0.426
0.430
0.455
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.008 (T7)
0.011 (T5/6)
0.009 (T6)
0.011 (T5/6)
0.020 (T7/8)
0.015 (T6/7)
0.011 (T6)
0.005 (T7)
0.004 (T7)
0.002 (T6)
0.007 (T5)
0.010 (>T7)
0.012 (T7)
0.008 (T6/7)
0.005 (T6/7)
0.006 (T5)
0.005 (>T7)
0.003 (T7)
0.007 (T5)
0.004 (T7)
0.004 (T5/6)
0.004 (T4)
0.008 (T5/6)
0.016 (T5)
0.003 (T7)
0.023 (T2/3)
0.004 (>T7)
0.003 (T5)
0.004 (T6)
0.012 (T7)
0.002 (>T7)
0.009 (T7/8)
0.004 (>T7)
0.003 (T6)
0.013 (T7)
0.037 (T4/5)
0.002 (T6)
0.006 (>T7)
0.016 (T6/7)
0.008 (T4)
0.019 (T5)
0.020 (T4/5)
CH4 -J
H2 O-J
0.398
0.451
0.459
0.497
0.301
0.386
0.453
0.374
0.394
0.481
0.540
0.231
0.379
0.443
0.410
0.467
0.304
0.399
0.533
0.412
0.437
0.574
0.451
0.509
0.374
0.726
0.337
0.517
0.448
0.375
0.303
0.333
0.320
0.436
0.390
0.569
0.448
0.260
0.379
0.611
0.426
0.460
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.015 (T6/7)
0.021 (<T7)
0.009 (<T7)
0.014 (<T7)
0.009 (T8)
0.008 (T7)
0.005 (<T7)
0.005 (T7)
0.003 (T7)
0.003 (<T7)
0.006 (<T7)
0.012 (T9)
0.011 (T7)
0.005 (<T7)
0.004 (<T7)
0.005 (<T7)
0.003 (T8)
0.002 (T6/7)
0.008 (<T7)
0.002 (<T7)
0.003 (<T7)
0.004 (<T7)
0.009 (<T7)
0.015 (<T7)
0.002 (T7)
0.023 (<T7)
0.003 (T8)
0.003 (<T7)
0.003 (<T7)
0.010 (T7)
0.001 (T8)
0.007 (T8)
0.004 (T8)
0.002 (<T7)
0.046 (T7/8)
0.071 (<T7)
0.001 (<T7)
0.005 (T9)
0.013 (T7)
0.006 (<T7)
0.015 (<T7)
0.016 (<T7)
WJ
±
±
±
±
±
±
±
±
±
0.172
0.261
0.295
0.236
0.318
0.202
0.263
0.189
0.327
0.156
0.314
0.346
0.282
0.191
0.144
0.203
0.290
0.295
0.343
0.295
0.168
±
±
±
±
±
±
±
±
±
±
±
±
0.392 ±
0.313 ±
0.243 ±
±
±
±
±
±
±
±
±
±
±
0.320
0.264
0.298
0.275
0.163
0.257
0.294
0.273
0.271
0.337
0.018 (T5/6)
0.016 (T6/7)
0.021 (T6)
0.010 (T6)
0.011 (T8)
0.030 (T6/7)
0.008 (T6)
0.010 (T6)
0.009 (T6)
0.003 (T5)
–
0.022 (T8)
0.023 (T6/7)
0.012 (T6)
0.017 (T7)
0.012 (T5/6)
0.007 (T7/8)
0.004 (T6/7)
0.015 (T7/8)
0.005 (T5)
–
–
0.045 (T4/5)
0.032 (T5/6)
0.004 (T7)
–
0.005 (T8)
0.005 (T6)
0.006 (T5)
0.021 (T6)
0.003 (T8)
0.018 (T8/9)
0.009 (T7/8)
0.005 (T6)
0.028 (T5/6)
0.033 (T5/6)
0.003 (T6)
0.011 (T8)
–
–
–
–
H2 O-H
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.187
0.407
0.292
0.266
0.117
0.263
0.173
0.301
0.267
0.553
0.287
0.134
±
±
±
±
±
±
±
±
±
±
±
±
0.309 ±
0.334 ±
0.204 ±
0.069
0.183
0.058
0.273
0.274
0.097
0.229
0.464
0.050
0.233
0.280
0.392
0.291
0.113
0.209
0.237
0.219
0.229
0.251
0.007 (T7)
0.011 (T6)
0.014 (T5)
0.009 (T6)
0.006 (>T7)
0.010 (T7)
0.008 (T7)
0.008 (T7)
0.008 (T7)
0.002 (T6/7)
–
0.012 (>T7)
0.016 (T7)
0.025 (>T7)
0.016 (T6)
0.008 (T6)
0.009 (>T7)
0.003 (T7)
0.011 (T5)
0.017 (>T7)
–
–
0.019 (T6)
0.022 (T6)
0.004 (T7)
–
0.003 (T7)
0.005 (T5)
0.005 (T6)
0.014 (T6)
0.002 (>T7)
0.013 (T6)
0.006 (T7)
0.004 (T6)
0.012 (T6)
0.018 (T4)
0.002 (T6)
0.008 (>T7)
–
–
–
–
CH4 -H
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Name
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.168
0.209
0.113
0.103
0.071
0.279
0.066
0.095
0.130
0.259
0.158
0.052
±
±
±
±
±
±
±
±
±
±
±
±
0.266 ±
0.166 ±
0.092 ±
0.260
0.090
0.175
0.091
0.152
0.102 ±
0.126 ±
0.118
0.127
0.172
0.100
0.030
0.084
0.079
0.108
0.056
0.027
0.010 (>T6)
0.012 (T6/7)
0.023 (T5/6)
0.012 (>T6)
0.027 (>T6)
0.019 (>T6)
0.013 (>T6)
0.010 (>T6)
0.025 (>T6)
0.012 (>T6)
–
0.023 (>T6)
0.017 (T6/7)
–
0.035 (T4/5)
0.008 (>T6)
0.022 (T5/6)
0.004 (>T6)
0.011 (T6)
–
–
–
0.025 (T4)
0.044 (T6/7)
0.014 (>T6)
–
0.014 (T5/6)
0.009 (T5)
0.012 (>T6)
0.015 (>T6)
0.010 (>T6)
0.028 (T4)
0.009 (>T6)
0.013 (>T6)
0.035 (T6/7)
0.068 (T4/5)
0.009 (T6)
0.011 (>T6)
–
–
–
–
CH4 -K
T5.5 S12
T6 S12
T8.5
T5.5 S12
T6.5 K11
T5.5 (photometric) S10
T7 (photometric) S10
T8.5
T8.5 A11
Note
Table 2. Spectral typing ratios for the confirmed T dwarfs as set out by Burgasser et al. (2006) and Burningham et al. (2009), along with the types from by-eye comparison to template spectral standards and the
final adopted types. The notes column indicates spectral types determined by authors where: A11 = Albert et al. (2011); K11 = Kirkpatrick et al. (2011); S10 = Scholz (2010); S12 = Scholz et al. (2012). In the
case of T8+ objects, the notes column indicates the spectral type using the Cushing et al. (2011) system.
470
B. Burningham et al.
Adopted
T6
T5
T5
T4.5
T6
T7.5
T7
T5
T5
T4.5
T6.5
T5.5
T6.5
T8p
T5
T5
T5
T6.5
T7
T6
T6
T5
T5
T3
T6
T6.5p
T4.5
T7
T7.5
T8
T4
T6
T6.5
T5.5p
Name
ULAS J1223−0131
ULAS J1258+0307
ULAS J1259+2933
ULAS J1302+1434
ULAS J1335+1506
ULAS J1338−0142
ULAS J1339−0056
ULAS J1339+0104
ULAS J1417+1330
ULAS J1421+0136
ULAS J1425+0451
ULAS J1449+1147
ULAS J1516+0110
WISE J1517+0529
ULAS J1534+0556
ULAS J1536+0155
ULAS J1549+2621
ULAS J1601+2646
ULAS J1614+2442
ULAS J1617+2350
ULAS J1619+2358
ULAS J1619+3007
ULAS J1626+2524
ULAS J1639+3232
ULAS J2116−0101
ULAS J2237+0642
ULAS J2300+0703
ULAS J2315+0344
ULAS J2318+0433
ULAS J2326+0201
ULAS J2331+0426
ULAS J2342+0856
ULAS J2352+1244
ULAS J2357+0132
T6
T5
T5
T4.5
T6
T7.5
T7
T5
T5
T4.5
T6.5
T5.5
T6.5
T8p
T5
T5
T5
T6.5
T7
T6
T7
T5
T5
T3
T6
T6.5p
T4.5
T7
T7.5
T8
T4
T6
T6.5
T5.5p
Templ.
0.147
0.187
0.227
0.276
0.148
0.010
0.073
0.215
0.204
0.374
0.122
0.235
0.126
0.045
0.256
0.298
0.233
0.084
0.096
0.181
0.074
0.321
0.243
0.387
0.175
0.128
0.316
0.093
0.054
0.052
0.388
0.156
0.152
0.173
0.037 (T6/7)
0.004 (T5)
0.004 (T5)
0.076 (T5/6)
0.007 (T5)
0.045 (T7/8)
0.003 (T7)
0.003 (T5)
0.003 (T5)
0.018 (T2)
0.008 (T6/7)
0.003 (T5)
0.004 (T6)
0.003 (T7)
0.015 (T5)
0.003 (T4)
0.004 (T5)
0.004 (T7)
0.010 (T6/7)
0.003 (T5)
0.018 (T6/7)
0.005 (T5)
0.014 (T4)
0.007 (T2/3)
0.011 (T5)
0.005 (T6)
0.009 (T4)
0.023 (T6/7)
0.005 (T7)
0.012 (>T7)
0.010 (T4)
0.003 (T6)
0.006 (T6)
0.004 (T5/6)
CH4 -J
0.315
0.383
0.408
0.432
0.371
0.237
0.275
0.402
0.381
0.610
0.284
0.378
0.323
0.220
0.395
0.475
0.381
0.224
0.289
0.383
0.273
0.441
0.468
0.584
0.377
0.295
0.504
0.270
0.248
0.169
0.489
0.340
0.312
0.362
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.019 (T6/7)
0.006 (T5)
0.005 (T5)
0.049 (T4/5)
0.010 (T6)
0.045 (>T7)
0.004 (T7/8)
0.004 (T5)
0.004 (T5)
0.015 (T3/4)
0.009 (T6/7)
0.004 (T5)
0.007 (T6/7)
0.004 (>T7)
0.016 (T5)
0.004 (T5)
0.005 (T5)
0.005 (T7)
0.012 (T7)
0.006 (T5/6)
0.017 (T7/8)
0.007 (T4/5)
0.020 (T5)
0.005 (T3)
0.010 (T5/6)
0.006 (T6/7)
0.008 (T4/5)
0.017 (T7)
0.007 (>T7)
0.006 (>T7)
0.013 (T3/4)
0.003 (T6)
0.008 (T6)
0.003 (T6)
H2 O-J
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.466
0.462
0.530
0.561
0.450
0.237
0.354
0.500
0.491
0.582
0.424
0.522
0.437
0.331
0.515
0.570
0.543
0.310
0.375
0.522
0.362
0.605
0.472
0.631
0.511
0.430
0.583
0.382
0.343
0.304
0.693
0.437
0.449
0.484
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.055 (<T7)
0.004 (<T7)
0.004 (<T7)
0.098 (<T7)
0.007 (<T7)
0.033 (T9)
0.003 (T7)
0.004 (<T7)
0.002 (<T7)
0.014 (<T7)
0.010 (<T7)
0.003 (<T7)
0.005 (<T7)
0.004 (T8)
0.016 (<T7)
0.003 (<T7)
0.004 (<T7)
0.004 (T8)
0.011 (T7)
0.004 (<T7)
0.012 (T7)
0.006 (<T7)
0.014 (<T7)
0.005 (<T7)
0.014 (<T7)
0.007 (<T7)
0.006 (<T7)
0.018 (T6/7)
0.006 (T8)
0.006 (T8)
0.013 (<T7)
0.002 (<T7)
0.009 (<T7)
0.005 (<T7)
WJ
±
±
±
±
±
±
±
±
±
±
±
±
±
±
0.276 ±
0.313 ±
0.184
0.257
0.203
0.466
0.291 ±
0.190 ±
0.389 ±
0.294 ±
0.263
0.301
0.301
0.224
0.351
0.348
0.318
0.278
0.062
0.313
0.027 (T5/6)
0.008 (T5)
0.009 (T4)
0.028 (T4/5)
0.015 (T6)
–
0.006 (T7)
0.006 (T6)
0.003 (T6)
–
0.016 (T6/7)
0.009 (T6)
0.009 (T6)
0.010 (T7)
0.025 (T4/5)
0.004 (T5)
0.007 (T5/6)
0.012 (T6)
0.030 (T9)
0.010 (T5/6)
–
0.015 (T4)
0.023 (T6)
–
0.014 (T6)
0.014 (T7/8)
–
0.024 (T7/8)
0.015 (T6/7)
0.011 (T7/8)
0.019 (T2/3)
–
0.014 (T6)
0.007 (T6)
H2 O-H
±
±
±
±
±
0.226 ±
0.288 ±
0.306 ±
0.312
0.332
0.382
0.378
0.289
CH4 -K
0.293 ± 0.058 (T4/5)
–
0.193 ± 0.007 (T5)
0.317 ± 0.054 (T3/4)
–
–
–
–
0.174 ± 0.002 (T6)
–
0.097 ± 0.021 (>T6)
0.254 ± 0.010 (T4)
–
0.201 ± 0.018 (T5)
0.237 ± 0.021 (T4/5)
–
–
0.080 ± 0.008 (>T6)
0.088 ± 0.023 (>T6)
0.170 ± 0.007 (T6)
–
0.199 ± 0.010 (T5)
–
–
0.138 ± 0.013 (T6/7)
0.139 ± 0.016 (T6/7)
–
0.254 ± 0.046 (T4/5)
0.077 ± 0.019 (>T6)
0.046 ± 0.010 (>T6)
0.429 ± 0.017 (T3)
–
0.091 ± 0.021 (>T6)
0.186 ± 0.004 (T5)
CH4 -H
0.434 ± 0.018 (T5)
–
0.415 ± 0.006 (T5)
0.516 ± 0.018 (T4)
–
–
–
–
0.378 ± 0.002 (T5)
–
0.282 ± 0.012 (T6)
0.389 ± 0.005 (T5)
–
0.168 ± 0.006 (T7)
0.442 ± 0.022 (T5)
–
–
0.205 ± 0.009 (T7)
0.160 ± 0.014 (T7/8)
0.348 ± 0.006 (T6)
–
0.402 ± 0.010 (T5)
–
–
0.339 ± 0.010 (T6)
0.241 ± 0.010 (T6/7)
–
0.193 ± 0.016 (T7)
0.182 ± 0.011 (T7)
0.097 ± 0.006 (>T7)
0.654 ± 0.016 (T3)
–
0.206 ± 0.014 (T7)
0.420 ± 0.007 (T5)
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Table 2 – continued
T7 (photometric) S10
T8 M13
T7 (photometric) S10
T5.5 (photometric) S10
Note
76 T dwarfs
471
472
B. Burningham et al.
Generally the sky levels were determined from annular regions.
Aperture corrections were taken from the IRAC handbook.
Table 4 lists the targets observed, the associated programme ID,
the date of the observation, the pipeline version used, the frame time,
the total integration time, the aperture size, the derived photometry
and uncertainties, and any notes on the data set. The table includes
known sources not presented in this work, for which the IRAC data
has not been previously published. We include it here so that it is
available to the community.
6 W I D E C O M M O N P RO P E R M OT I O N
B I N A RY S Y S T E M S
6.1 Identifying companions
5.2 Warm-Spitzer photometry
Warm-Spitzer IRAC photometry was obtained for some of the T
dwarfs presented in this work via Cycles 6, 7 and 8 GO programmes
60093, 70058 and 80077 (PI Leggett). The observations were carried out in both the [3.6] and [4.5] bands (hereafter Ch1 and Ch2,
respectively), with typically a 30 s frame time, repeated three to six
times per pointing, and dither patterns consisting of 12 or 16 positions. In all cases the post-basic-calibrated-data mosaics generated
by the Spitzer pipeline were used to obtain aperture photometry.
Table 3. WISE photometry for 25 of the spectroscopically confirmed T dwarfs presented here for the first time.
Name
CFBDS J0133+0231
ULAS J0139+1503
ULAS J0200+0908
ULAS J0329+0430
ULAS J0745+2332
ULAS J0758+2225
ULAS J0819+2103
ULAS J0821+2509
ULAS J0950+0117
ULAS J0954+0623
ULAS J1021+0544
ULAS J1029+0935
ULAS J1043+1048
ULAS J1152+0359
ULAS J1152+1134
ULAS J1204−0150
ULAS J1206+1018
ULAS J1338−0142
ULAS J1417+1330
ULAS J1449+1147
ULAS J1517+0529
ULAS J1549+2621
ULAS J1639+3232
ULAS J2326+0201
ULAS J2342+0856
W1
W2
W3
W4
WISE blend?
17.78 ± 0.27
17.86 ± 0.26
16.16 ± 0.07
17.52 ± 0.25
14.79 ± 0.04
16.95 ± 0.16
17.16 ± 0.16
17.40 ± 0.22
18.05 ± 0.34
16.67 ± 0.13
16.63 ± 0.14
16.84 ± 0.13
>18.28
16.97 ± 0.15
16.89 ± 0.15
16.66 ± 0.11
17.55 ± 0.24
>18.22
16.67 ± 0.08
17.39 ± 0.16
>18.10
17.13 ± 0.10
16.62 ± 0.10
18.03 ± 0.42
16.07 ± 0.08
15.10 ± 0.09
15.94 ± 0.16
15.70 ± 0.14
15.35 ± 0.14
14.47 ± 0.07
15.23 ± 0.12
15.28 ± 0.10
15.48 ± 0.14
14.48 ± 0.06
14.66 ± 0.08
15.31 ± 0.28
14.29 ± 0.08
15.66 ± 0.18
15.34 ± 0.13
14.66 ± 0.08
14.70 ± 0.08
15.83 ± 0.19
16.12 ± 0.19
14.70 ± 0.06
14.84 ± 0.07
15.13 ± 0.08
16.09 ± 0.13
15.02 ± 0.08
15.45 ± 0.16
13.97 ± 0.05
>12.91
>12.93
>12.88
>12.20
12.51 ± 0.50
>12.65
>12.17
>12.67
>12.85
>12.64
>11.94
11.58 ± 0.33
12.10 ± 0.32
12.29 ± 0.38
12.37 ± 0.41
12.48 ± 0.42
>12.36
>13.01
12.51 ± 0.29
>12.40
>13.13
13.42 ± 0.5
12.28 ± 0.28
>12.53
12.63 ± 0.53
>9.26
>9.31
>9.24
>8.53
>8.41
>9.03
>8.81
>9.17
>9.20
>8.70
>8.14
>8.58
>9.07
>9.03
>8.57
>8.60
>8.86
>9.34
>9.49
>9.44
>9.52
>9.89
>9.05
>8.96
9.08 ± 0.54
N
N
Y
N
Y
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
N
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 4. A Y − J versus J − H colour–colour diagram for 146 spectroscopically confirmed T dwarfs in the UKIDSS LAS with YJH photometry.
Spectral types are indicated by coloured symbols.
In Table 5, we present proper motions for the targets identified
in this paper, along with those for late-T dwarfs found within the
UKIDSS LAS sky in previous works. The vast majority of these
proper motions have been drawn from the catalogue of Smith et al.
(in preparation) which presents proper motions calculated from two
epochs of J-band UKIDSS LAS observations. Here, we present
only the absolute proper motions, and refer the reader to the Smith
et al. catalogue for further astrometric parameters. A small number
of additional proper motions have been calculated using an identical method to that used in Smith et al. (in preparation), but using
our follow-up WFCAM J-band observations for the second epoch,
instead of UKIDSS survey data.
We have cross-matched our full catalogue of T dwarfs identified within the UKIDSS LAS sky that have proper motions (128
targets) against several astrometric catalogues. We cross-matched
our targets against the Hipparcos (Perryman et al. 1997; van
Leeuwen 2007), LSPM-NORTH (Lépine & Shara 2005) and NOMAD (Zacharias et al. 2004) catalogues. We searched projected
separations up to 200 00 au, assuming a minimum likely distance
76 T dwarfs
473
Figure 6. A histogram of UKIDSS and WISE + UKIDSS detected T dwarfs
in our sample for objects with 18.3 < J < 18.8.
for each source. Minimum and maximum likely distances for each
source were determined by considering the ±0.5 subtype spectral
type uncertainty, the mean MJ for each spectral subtype and the
scatter about MJ as presented in Dupuy & Liu (2012). Since a significant fraction of wide-common proper motion companions are
themselves multiple systems (e.g. Faherty et al. 2010), we also
assigned an upper limit to the distance based on the target being
an unresolved equal luminosity binary system. These distances are
presented in Table 5.
To identify apparent pairs with common proper motion we selected only objects with total proper motions that are more than
3σ significant, and greater than 100 mas yr−1 (92 objects). To be
considered common proper motion pairs we required 4σ matches
in both μαcos δ and μδ .
To assess if possible pairs share a common distance (in the absence of trigonometric parallax), we estimated the maximum and
minimum plausible absolute magnitudes of the candidate primary
stars based on the hypothesis that they lie at the same distance as
their candidate companions. The candidate primary stars are compared to Hipparcos stars with VJ photometry in Fig. 7. Those targets
whose maximum and minimum hypothesised MJ bracket the main
sequence (or white dwarf or giant branches) were accepted as candidate common proper motion binary pairs to our T dwarfs. Of our
nine candidate primaries with 4σ matched proper motions and VJ
colours, five appear very likely to have common distance to our T
dwarfs. One more has a minimum hypothetical MJ value that lies on
the periphery of the main sequence and is thus consistent with the
target sharing a common distance to the T dwarf, if the T dwarf is
itself an unresolved binary. Three pairs are ruled out by the common
distance test. The initial characterization of these pairs is given in
Table 6.
It is worth highlighting that this method restricts us to investigating stars with V- and J-band photometry in the NOMAD and LSPM
catalogues, and we have likely thus excluded a number of genuine binary companions. For example, we also recovered the white
dwarf–T dwarf pair LSPM 1459+0851AB (Day-Jones et al. 2011)
as a proper motion match; however, the lack of 2MASS photometry
for the WD primary excluded it from our analysis at this stage. This
was the only LSPM candidate with proper motion agreement that
lacked VJ photometry. However, a large number of NOMAD candidate primaries with proper motion agreement lacked appropriate
photometry for the common distance check. It is thus likely that a
number of additional binary companions may remain unidentified
in our sample, particularly for more distant red primaries that lack
Tycho photometry (e.g. M dwarfs).
Of the five strong candidates, three are previously identified binary systems in our sample: Ross 458ABC (Goldman et al. 2010;
Scholz 2010), BD+01 2930AB (Pinfield et al. 2012) and Hip
73786AB (LHS 3020AB Scholz 2010; Murray et al. 2011) and
two are new candidates: ULAS J0950+0117 (T8) + LHS 6176
(estimated M4); ULAS J1339+0104 (T5) + HD 118865 (F8). The
latter of these was also identified in our cross-match against Hipparcos, and the parallax for the primary is consistent with our estimated
distance to the T dwarf secondary. The former has been independently identified as a candidate proper motion pair by Luhman et al.
(2012) since our detailed study of it had already commenced.
To assess the probability of chance alignment for our new candidate binary pairs, we followed the method described in Dhital et al.
(2010), which calculates the frequency of unrelated pairings using a Galactic model that is parametrized by empirically measured
stellar number density (Jurić et al. 2008; Bochanski et al. 2010)
and space velocity (Bochanski et al. 2007) distributions. All stars
in the model are single (and hence unrelated); therefore any stars
within the 5D ellipsoid defined by the binary’s position, angular
separation, distance and proper motions is a chance alignment. We
performed 106 Monte Carlo realizations to calculate the probability
of chance alignment. The chance-alignment probabilities for the
three new candidates are given in the ‘Notes’ column of Table 6.
The weaker candidate has a correspondingly higher probability of
chance alignment, and in the absence of further data on the primary star, or improved distance estimates, it is not reasonable to
pursue further analysis of this candidate system. The two strongest
new candidates, however, are likely to be bona fide common proper
motion systems and we proceed on this basis.
6.2 LHS 6176AB
6.2.1 Distance to LHS 6176AB
As part of our wider campaign for determining accurate distances to
late-T dwarfs in our sample (e.g. Marocco et al. 2010; Smart et al.
2010), we have obtained a trigonometric parallax measurement for
ULAS J0950+0117 (LHS 6176B), and also for the proposed primary LHS 6176. The astrometric observations and image reduction
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 5. A histogram of UKIDSS and WISE + UKIDSS detected T dwarfs
in our sample for objects with J < 18.3.
474
B. Burningham et al.
Table 4. Spitzer photometry for a subset of T dwarfs, selected either as examples of late-T types, peculiar spectra or as benchmark objects. Full designations
are given for those objects whose discoveries are reported in other publications, whilst abbreviated names are given for those objects whose discovery is
presented in this work. Notes: A: The data were taken when Spitzer was cold, and longer wavelength photometry was also obtained: [5.8] = 14.24 ± 0.03,
[8.0] = 13.31 ± 0.03. B: Separate, non-annular, skies used due to background structure.
Name
Spectral
type
Programme
number
Obs date
(UT)
Pipeline
version
Frame
time (s)
Integration
(min)
Aperture
(arcsec)
[3.6]
(mag)
[4.5]
(mag)
Notes
CFBDS J005910.90−011401.3
CFBDS J030135−161418
2MASS J07290002−3954043
ULAS J0809+2126
CFBDS J092250+152741
ULAS J0950+0117
ULAS J1043+1048
ULAS J123327.45+121952.2
ULAS J1339+0104
ULAS J2237+0642
ULAS J2326+0201
T8.5a
T7.5b
T8c
T8
50667
60093
60093
70058
60093
60093
70058
80077
80077
80077
80077
2009-01-26
2009-09-02
2009-12-04
2011-05-31
2010-05-29
2010-01-06
2012-02-01
2012-09-03
2012-09-18
2012-09-08
2012-01-31
S18.7.0
S18.12.0
S18.13.0
S18.18.0
S18.18.0
S18.13.0
S19.1.0
S19.1.0
S19.1.0
S19.1.0
S19.1.0
12
30
12
30
30
30
30
30
30
30
30
7.2
48.0
1.0
48.0
48.0
48.0
48.0
48.0
48.0
48.0
24.0
7.2
7.2
4.8
4.8
7.2
7.2
7.2
7.2
7.2
7.2
7.2
15.71 ± 0.01
16.95 ± 0.01
14.47 ± 0.01
17.74 ± 0.05
17.45 ± 0.04
16.28 ± 0.01
16.94 ± 0.01
16.83 ± 0.01
16.93 ± 0.01
17.87 ± 0.01
16.84 ± 0.03
13.66 ± 0.01
15.42 ± 0.01
12.95 ± 0.01
16.03 ± 0.01
16.11 ± 0.01
14.35 ± 0.02
15.34 ± 0.03
15.61 ± 0.01
16.08 ± 0.02
15.48 ± 0.02
15.37 ± 0.01
A
B
et al. (2011); b Albert et al. (2011); c Looper, Kirkpatrick & Burgasser (2007); d Burningham et al. (2010b).
procedures were identical to those described in Smart et al. (2010).
Since the observing strategy was optimized for measuring the distance to the T dwarf, LHS 6176A was often close to saturation on
the image and, as a result, the centroiding precision is reduced. For
this reason, the original survey image from 2008 could not be used
at all in the solution, so the final astrometric parameter precision for
LHS 6176 is much lower than for the T dwarf. In total, 35 observations with 78 reference stars over a baseline of 4.22 yr were used
for LHS 6176B, compared to 33 observations with 525 reference
stars over 2.16 yr for LHS 6176A. The greater number of reference
stars for LHS 6176A is due to their being drawn from the entire
WFCAM chip, rather than from just the immediate vicinity of the
target. All astrometric parameters indicate a common distance and
common motion for the two objects, supporting our interpretation
of the pair as a binary system. The proper motions are also consistent with those found for LHS 6176 in Lépine & Shara (2005)
(249, −346 mas yr−1 ). The resulting astrometric parameters for the
pair are given in Table 7. When we re-run our chance alignment estimate calculation using the new trigonometric distance estimates,
the probability of chance alignment is smaller than one part in 107 .
Assuming the more precise distance to LHS 6176B as the distance
to the system, we find that the projected separation of the pair is
thus 970 au.
6.2.2 Spectroscopy of LHS 6176A
Optical spectroscopy of LHS 6176A was obtained on the night
of 2012 May 5 with the SuperNova Integral Field Spectrograph
(SNIFS; Lantz et al. 2004) on the University of Hawaii 2.2 m
telescope on Mauna Kea. SNIFS was operated with a dichroic mirror
that separated the incoming light into blue (3200–5200 Å) and red
(5100–9700 Å) spectrograph channels as well as an imaging channel
that was used for guiding. The observations yielded a resolution of
R ≈ 1000 for the blue channel and R ≈ 1300 for the red. Integration
time was 210 s, which was sufficient for high S/N (≈80 per Å) in
the red channel, which is our region of interest for spectral typing
an object as red as LHS 6176A. Basic reduction was performed
automatically by the SNIFS processing pipeline, which included
dark, bias, flat-field corrections, and cleaning of bad pixels and
cosmic rays. The clean data were then assembled into blue and red
data cubes. Wavelengths were calibrated with arc lamp exposures
taken at the same telescope pointing as the science data (to correct
for flexures).
The calibrated spectrum was then sky-subtracted, and a 1D spectrum was extracted from the image cube using a point spread function model. Observations of the Feige 66, BD+75325 and G191B2B
spectrophotometric standards (Oke 1990) taken over the course of
the night were used to correct each spectrum for instrument response
and remove telluric lines. Spectra were then shifted in wavelength to
the rest frames of their source stars by cross correlating each spectrum with similar spectral type templates from the SDSS (Stoughton
et al. 2002; Bochanski et al. 2007). More details on the SNIFS data
processing pipeline can be found in Bacon et al. (2001) and Aldering
et al. (2006), and more information on additional data processing
can be found in Lepine et al. (2013).
Fig. 8 shows the extracted red channel. Overplotted is an M4
template spectrum based on mean spectra of inactive M dwarfs from
Bochanski et al. (2007), their extremely close agreement across
nearly the entire wavelength range covered by our data leads us
to adopt a spectral type of M4V for LHS 6176. Their excellent
agreement, and the lack of Hα in emission also suggests a lack of
activity for this object. The absence of Hα in emission has been
found to be typical of M4 dwarfs with ages of >3.5 Gyr (West et al.
2008).
We have also obtained NIR spectroscopy of LHS 6176A on 2012
April 30 UT at the NASA Infrared Telescope Facilty on the summit
of Mauna Kea, Hawaii. We used the facility spectrograph SpeX
(Vacca et al. 2003) in short-wavelength cross-dispersed mode with
the 0.3 arcsec slit, which provided an average spectral resolution
(R ≡ λ/λ) of ≈2000. We obtained six exposures, each with a
120 s integration time and dithered in an ABBA pattern, for a total
of 12 min on-source. We observed the A0 V star HD 92245 contemporaneously for telluric calibration. All spectra were reduced
using version 3.4 of the SPEXTOOL software package (Vacca et al.
2003; Cushing et al. 2004). The resulting JHK spectrum is shown
in Fig. 9.
We have taken three approaches to estimating the metallicity of
LHS 6176A
Method 1. Our parallax for LHS 6176 also allows us to estimate
its metallicity using the improved [Fe/H] versus MKs /V − Ks calibrations of Schlaufman & Laughlin (2010) and Neves et al. (2012).
Since the uncertainty on the parallax of the T dwarf component is
considerably smaller than that of LHS 6176A, we have adopted the
former’s distance for the system. To maximize the precision of our
metallicity estimate we have obtained new V-band photometry of
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
a Cushing
T8p
T8
T3.5pd
T5
T6.5p
T8
B
α
(J2000)
01:30:17.79
02:00:13.18
02:26:03.18
02:45:57.88
03:25:53.13
03:29:20.22
07:41:49.01
07:45:03.03
07:46:17.17
07:45:02.79
07:50:03.84
07:55:47.95
07:58:29.83
07:59:37.75
08:09:18.41
08:11:10.86
08:14:07.51
08:15:07.26
08:19:18.58
08:19:48.08
08:21:55.49
08:23:27.46
08:27:07.67
08:30:48.89
08:37:56.19
08:40:36.72
08:42:11.68
08:51:39.03
08:53:42.94
09:01:16.23
09:13:09.55
09:26:24.76
09:29:06.75
09:29:26.44
09:32:45.48
09:38:29.28
09:39:51.04
09:43:31.49
09:43:49.60
09:45:16.39
09:48:06.06
09:50:47.28
ULAS J013017.79+080453.9
ULAS J020013.18+090835.2
ULAS J022603.18+070231.4
ULAS J024557.88+065359.4
SDSS J032553.17+042540.1
ULAS J032920.22+043024.5
SDSS J074149.14+235127.3
ULAS J074503.03+233240.3
ULAS J074617.17+235532.2
ULAS J074502.79+245516.3
WISEP J075003.84+272544.8
2MASS J07554795+2212169
ULAS J075829.83+222526.7
ULAS J075937.75+185555.0
ULAS J080918.41+212615.2
ULAS J081110.86+252931.8
ULAS J081407.51+245200.9
ULAS J081507.26+271119.2
ULAS J081918.58+210310.4
ULAS J081948.08+073323.3
ULAS J082155.49+250939.6
ULAS J082327.46+002424.4
ULAS J082707.67−020408.2
SDSS J083048.81+012831.0
ULAS J083756.19−004156.0
ULAS J084036.72+075933.6
ULAS J084211.68+093611.9
ULAS J085139.03+005340.9
ULAS J085342.94+000651.8
ULAS J090116.23−030635.0
ULAS J091309.55−003136.9
ULAS J092624.76+071140.7
WISEP J092906.77+040957.9
ULAS J092926.44+110547.3
ULAS J093245.48+310206.4
ULAS J093829.28−001112.6
ULAS J093951.0400:6:3.60
ULAS J094331.49+085849.2
ULAS J094349.60+094203.4
ULAS J094516.39+075545.6
ULAS J094806.06+064805.0
ULAS J095047.28+011734.3
Spectral type
T6
T6
T7
T7
T5.5
T5
T5
T9
T7
T6.5
T8.5
T5
T6.5
T6
T8
T7
T5p
T7p
T6
T6p
T4.5
T4.0
T5.5
T4.5
T3.0
T4.5
T5.5
T4.0
T6p
T7.5
T6
T3.5
T7
T2
T2
T4.5
T5.5
T5p
T4.5p
T5
T7
T8p
δ
(J2000)
+08:04:53.90
+09:08:35.20
+07:02:31.40
+06:53:59.40
+04:25:40.10
+04:30:24.50
+23:51:25.90
+23:32:40.30
+23:55:32.20
+24:55:16.30
+27:25:44.80
+22:12:14.50
+22:25:26.70
+18:55:55.00
+21:26:15.20
+25:29:31.80
+24:52:00.90
+27:11:19.20
+21:03:10.40
+07:33:23.30
+25:09:39.60
+00:24:24.40
−02:04:08.20
+01:28:28.90
−00:41:56.00
+07:59:33.60
+09:36:11.90
+00:53:40.90
+00:06:51.80
−03:06:35.00
−00:31:36.90
+07:11:40.70
+04:09:57.70
+11:05:47.30
+31:02:06.40
−00:11:12.60
+00:16:53.60
+08:58:49.20
+09:42:03.40
+07:55:45.60
+06:48:05.00
+01:17:34.30
6
6
5
5
5
7
5
3
5
6
7
4
6
5
5
5
5
7
5
3
4
2
1
Ref.
12.98
7.16
54.55
− 101.13
− 194.65
225.25
− 259.64
− 251.7
− 3.75
− 84.88
− 734.03
− 4.2
− 105.34
− 48.17
− 152.69
45.25
− 51.29
− 50.15
− 57.72
13.17
− 449.31
− 35.13
23.21
186.16
− 13.13
− 270.68
− 201.73
− 61.66
− 43.79
− 56.37
72.67
− 51.48
526.12
− 41.48
− 44.82
− 255.55
159.08
− 83.54
44.98
− 129.17
238.8
242.68
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
24.45
27.97
16.02
32.52
15.04
15.8
11
21.34
24.51
12.55
14.2
10.28
10.37
15.64
15.28
11.99
14.3
37.06
11.49
9.41
14.02
10.06
8.01
9.08
10.51
12.98
12.49
11.07
9.51
8.62
11.54
10.67
32.38
11.88
14.75
12.59
11.15
11.48
13.17
11.35
12.16
11.79
μαcos δ
(mas yr−1 )
− 43.98
− 40.39
59.24
− 112.72
− 102.68
− 59.84
− 216.24
− 287.98
− 120.29
− 60.28
− 195.27
− 252.3
− 57.49
− 81.2
− 154.7
− 231.89
− 9.48
− 79.91
− 181.4
− 68.5
− 56.45
− 221.22
− 111.73
− 361.52
− 94.16
− 82.66
− 53.25
− 39.68
120.23
− 253.84
− 51.54
− 420.41
− 438.45
9.95
− 6.29
− 81.28
− 299.34
− 78.78
− 123.92
− 41.17
− 273.6
− 386.56
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
59.09
28.01
16.38
29.84
14.86
15.51
11.02
21.1
24.53
12.52
14.49
10.65
10.84
15.91
16.28
12.05
14.11
36.98
10.67
9.4
10.26
10.11
7.97
8.96
10.53
12.93
12.62
10.86
9.62
10.39
11.36
12.3
30.96
11.8
14.51
11.48
11.03
11.5
12.72
10.61
11.96
11.71
μδ
(mas yr−1 )
Dmin
(pc)
30.1
28.4
22.6
21.0
13.4
31.5
14.6
8.7
28.2
31.6
12.9
12.1
24.5
42.9
12.0
14.6
49.7
20.5
19.2
34.7
31.3
55.4
22.7
18.2
48.8
72.1
39.8
62.5
41.5
16.1
50.4
34.5
10.6
63.1
53.7
57.3
31.6
51.1
65.8
30.6
26.0
9.3
Baseline
(yr)
1.29W
1.28
2.02W
1.05W
2.09W
2.09W
2.03
2.03
2.03
2.03
2.00
2.91
2.91
2.9
1.86
1.97
2.03
0.90W
1.86
5.09
2.03
5.04
6.07
5.07
6.06
5.28
5.20
5.18
5.19
6.18
5.28
5.12
0.97W
5.12
2.02
5.32
5.32
5.21
4.80
5.33
6.06
3.97
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Name
51.3
48.4
53.0
49.2
24.1
53.2
24.7
19.4
66.1
45.1
36.4
20.4
35.0
73.1
44.3
34.2
83.9
48.1
32.7
59.2
45.9
83.9
40.8
26.7
78.0
105.7
71.4
94.6
70.8
31.3
85.9
47.7
24.8
100.5
85.6
83.9
56.8
86.3
96.4
51.8
61.0
34.2
Dmax
(pc)
72.4
68.4
74.8
69.5
34.0
75.2
34.9
27.4
93.3
63.7
51.4
28.9
49.4
103.3
62.5
48.3
118.6
67.9
46.2
83.6
64.9
118.5
57.6
37.7
110.2
149.3
100.9
133.7
100
44.2
121.3
67.4
35.0
141.9
120.9
118.6
80.2
121.9
136.1
73.1
86.2
48.3
Dmax (binary)
(pc)
Table 5. Proper motions for T dwarfs within UKIDSS LAS sky. Unless otherwise stated proper motions are drawn from the catalogue of Smith et al. (in preparation). Epoch baselines denoted with W indicate
that the proper motion has been calculated from our own WFCAM follow-up following the same method as used for the Smith et al. catalogue. Spectral types are on the system of Burningham et al. (2008).
Maximum and minimum plausible distances have been calculated using the mean magnitudes for each spectral subtype from Dupuy & Liu (2012) and assuming ±0.5 subtype precision on types. An additional
maximum distance to account for possible unresolved binaries is given in the final column.
76 T dwarfs
475
α
(J2000)
09:54:29.90
09:58:29.86
10:01:13.04
10:07:59.90
10:12:43.54
10:17:21.40
10:18:21.78
10:21:44.87
10:23:05.44
10:29:40.52
10:34:34.52
10:43:55.37
10:51:34.32
10:52:35.42
11:10:09.85
11:49:25.58
11:50:38.79
11:53:38.74
11:55:08.39
11:57:18.02
11:57:59.04
12:02:57.05
12:04:44.67
12:06:21.03
12:07:44.65
12:12:26.80
12:23:43.35
12:31:53.60
12:33:27.45
12:38:28.51
12:39:03.75
12:48:04.56
12:57:08.07
12:58:35.97
13:00:41.73
13:02:17.21
13:02:27.54
13:03:03.54
13:15:08.42
13:19:43.77
13:20:48.12
13:26:05.18
13:35:02.11
13:35:53.45
13:38:28.69
Name
ULAS J095429.90+062309.6
ULAS J095829.86−003932.0
CFBDS J100113+022622
ULAS J100759.90−010031.1
ULAS J101243.54+102101.7
ULAS J101721.40+011817.9
ULAS J101821.78+072547.1
ULAS J102144.87+054446.1
ULAS J102305.44+044739.2
ULAS J102940.52+093514.6
ULAS J103434.52−001553.0
ULAS J104355.37+104803.4
ULAS J105134.32−015449.8
ULAS J105235.42+001632.7
SDSS J111010.01+011613.1
ULAS J114925.58−014343.2
ULAS J115038.79+094942.9
ULAS J115338.74−014724.1
ULAS J115508.39+044502.3
ULAS J115718.02−013923.9
ULAS J115759.04+092200.7
ULAS J120257.05+090158.8
ULAS J120444.67−015034.9
ULAS J120621.03+101802.9
ULAS J120744.65+133902.7
ULAS J121226.80+101007.4
ULAS J122343.35−013100.7
ULAS J123153.60+091205.4
ULAS J123327.45+121952.2
ULAS J123828.51+095351.3
ULAS J123903.75+102518.6
ULAS J124804.56+075904.0
ULAS J125708.07+110850.4
ULAS J125835.97+030736.1
ULAS J130041.73+122114.7
ULAS J130217.21+130851.2
ULAS J130227.54+143428.0
ULAS J130303.54+001627.7
ULAS J131508.42+082627.4
ULAS J131943.77+120900.2
ULAS J132048.12+102910.6
ULAS J132605.18+120009.9
ULAS J133502.11+150653.5
ULAS J133553.45+113005.2
ULAS J133828.69−014245.4
Spectral type
T5
T5.5
T5
T5.5
T5.5
T8p
T5
T6
T6.5
T8
T6.5p
T8
T6
T5
T5.5
T5
T6.5
T6
T7
T5
T2.5
T5
T4.5
T5
T6.0
T5
T6
T4.5p
T4p
T8.5
T0
T7
T4.5
T5
T8.5
T8.5
T4.5
T5.5
T7.5
T5
T5
T6p
T6
T9
T7.5
δ
(J2000)
+06:23:09.60
−00:39:32.00
+02:26:22.40
−01:00:31.10
+10:21:01.70
+01:18:17.90
+07:25:47.10
+05:44:46.10
+04:47:39.20
+09:35:14.60
−00:15:53.00
+10:48:03.40
−01:54:49.80
+00:16:32.70
+01:16:10.50
−01:43:43.20
+09:49:42.90
−01:47:24.10
+04:45:02.30
−01:39:23.90
+09:22:00.70
+09:01:58.80
−01:50:34.90
+10:18:02.90
+13:39:02.70
+10:10:07.40
−01:31:00.70
+09:12:05.40
+12:19:52.20
+09:53:51.30
+10:25:18.60
+07:59:04.00
+11:08:50.40
+03:07:36.10
+12:21:14.70
+13:08:51.20
+14:34:28.00
+00:16:27.70
+08:26:27.40
+12:09:00.20
+10:29:10.60
+12:00:09.90
+15:06:53.50
+11:30:05.20
−01:42:45.40
9
6
6
5
5
5
10
5
5
5
9
5
5
5
5
5
6
5
5
4
5
6
5
5
6
8
7
5
9
7
Ref.
− 494.26
− 59.18
− 90.21
− 226.6
− 390.6
− 83
− 168.38
− 29.92
15.49
− 407.57
− 100.45
96.75
− 65.76
− 30.88
− 235.42
− 112.51
− 94.21
− 568.28
483.25
108.24
− 90.15
− 42.73
− 402.6
− 400.83
− 154.53
− 181.77
− 168.28
82.62
35.8
− 450.71
− 209.49
− 230.78
28.65
− 170.86
− 635.95
− 427.96
− 42.14
12.3
− 36.52
− 121.9
102.21
70.63
3.02
− 183.38
200.75
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
9.92
12.49
13.7
11.08
11.94
12.46
12.88
36.78
36.22
37.19
14.95
37.4
32.13
18.54
34.43
13.11
13.94
16.79
14.18
13.62
11.89
12.09
15.73
19.5
12.49
16.23
15.18
15.08
10.74
15.46
13.53
12.95
14.42
14.19
14.43
13.41
16.29
20.99
12.73
16.04
11.03
11.04
14.16
13.1
21.02
μαcos δ
(mas yr−1 )
− 436.26
1.73
49.41
145.69
− 555
− 15.14
− 14.81
30.79
− 83.08
− 145.97
− 30.41
− 73.04
− 40.03
− 140.23
− 273.03
11
− 19.38
− 333.86
− 533.47
− 20.85
30.25
− 62.13
132.24
− 87.42
0.84
− 123.53
65.79
− 80.16
92.94
41.72
92.23
− 145.35
− 57.13
− 30.75
− 27.9
− 9.34
− 14.66
− 274.16
− 103.85
− 22.9
− 56.56
− 29.76
− 106.03
− 214.17
− 109.89
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
10.79
12.33
14.23
11.19
15.99
11.74
10.33
32.49
37.91
26.38
14.54
37.95
32.4
18.4
31.09
12.55
14.08
13.12
13.2
12.8
11.08
10.63
11.38
18.75
12.29
16.5
14.32
13.8
11.15
14.17
13.62
11.95
13.94
12.34
12.29
12.92
15.2
19.97
12.22
14.59
11.44
12.02
12.56
10.92
20.2
μδ
(mas yr−1 )
Dmin
(pc)
20.3
51.8
59.1
45.7
19.9
11.7
33.8
26.5
35.0
6.6
43.5
10.2
27.7
56.8
14.2
40.7
40.0
25.7
20.7
42.1
25.4
21.4
25
64.6
35.3
53.2
42.8
71.4
40.7
12.8
65.2
15.6
53.5
46.1
5.1
9.7
58.9
53.5
25.2
58.6
35.6
24.7
30.6
5.5
23.3
Baseline
(yr)
2.75
6.24
3.99
6.08
4.74
5.05
6.00
0.90W
0.91W
0.91W
4.01
0.91W
0.99W
2.39
0.95W
4.05
5.59
3.90
2.96
3.90
4.79
4.79
3.64
3.85
4.79
3.86
3.64
4.96
4.89
4.96
4.96
5.03
4.81
2.83
3.83
4.88
3.71
2.64
5.67
4.87
4.94
4.87
3.85
4.86
2.56
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Table 5 – continued
34.4
92.9
99.8
82.1
35.6
43.3
57.2
45.3
49.9
24.3
61.9
37.7
47.2
95.9
25.5
68.9
57.0
43.9
48.5
71.1
34.4
36.1
36.6
109.1
60.3
89.9
73.0
104.7
61.7
36.1
86.3
36.6
78.3
78.0
14.3
27.3
86.3
95.9
49.0
99.1
60.3
42.1
52.2
12.1
45.3
Dmax
(pc)
48.5
131.2
141.0
115.9
50.4
61.1
80.7
64.0
70.5
34.4
87.5
53.2
66.7
135.5
36.0
97.3
80.5
61.9
68.6
100.5
48.5
51.1
51.8
154.2
85.1
127.1
103.1
147.9
87.1
51.1
121.9
51.7
110.7
110.2
20.2
38.5
121.9
135.5
69.2
140.0
85.1
59.4
73.8
17.1
64.0
Dmax (binary)
(pc)
476
B. Burningham et al.
13:39:33.64
13:39:43.79
13:46:46.10
13:49:40.81
13:56:07.41
14:16:23.94
14:17:56.22
14:21:45.63
14:23:20.79
14:25:36.35
14:44:58.87
14:45:55.24
14:49:01.91
14:52:43.59
14:59:35.25
15:01:35.33
15:04:11.73
15:04:57.66
15:05:47.89
15:16:37.89
15:17:21.12
15:25:26.25
15:26:55.78
15:29:12.23
15:34:06.06
15:36:53.80
15:44:27.34
15:47:01.84
15:49:14.45
16:01:43.75
16:14:36.96
16:17:10.39
16:19:34.78
16:19:38.12
16:26:55.04
16:39:31.52
22:37:28.91
23:00:49.08
23:42:28.97
23:57:15.98
ULAS J133933.64−005621.1
ULAS J133943.79+010436.4
SDSS J134646.43−003150.3
ULAS J134940.81+091833.3
ULAS J135607.41+085345.2
ULAS J141623.94+134836.3
ULAS J141756.2213:0:5.80
ULAS J142145.63+013619.0
ULAS J142320.79+011638.2
ULAS J142536.35+045132.3
ULAS J144458.87+105531.1
ULAS J144555.24+125735.1
ULAS J144901.91+114711.4
ULAS J145243.59+065542.9
ULAS J145935.25+085751.2
ULAS J150135.33+082215.2
SDSS J150411.63+102718.4
ULAS J150457.66+053800.8
ULAS J150547.89+070316.6
ULAS J151637.89+011050.1
WISE J151721.13+052929.3
ULAS J152526.25+095814.3
CFBDS J152655.78+034534.8
ULAS J152912.23+092228.5
ULAS J153406.06+055643.9
ULAS J153653.8001:5:0.60
ULAS J154427.34+081926.6
ULAS J154701.84+005320.3
ULAS J154914.45+262145.6
ULAS J160143.75+264623.4
ULAS J161436.96+244230.1
ULAS J161710.39+235031.4
ULAS J161934.78+235829.3
ULAS J161938.12+300756.4
ULAS J162655.04+252446.8
ULAS J163931.52+323212.7
ULAS J223728.91+064220.1
ULAS J230049.08+070338.0
ULAS J234228.97+085620.1
ULAS J235715.98+013240.3
Spectral type
T7
T5
T6.5
T7
T5
T7.5p
T5
T4.5
T8p
T6.5
T5
T6.5
T5.5
T4.5
T4.5
T4.5
T7
T6p
T4
T6.5
T8p
T6.5
T4
T6
T5
T5
T3.5
T5.5
T5
T6.5
T7
T6
T6
T5
T5
T3
T6.5p
T4.5
T6
T5.5p
δ
(J2000)
−00:56:21.10
+01:04:36.40
−00:31:51.40
+09:18:33.30
+08:53:45.20
+13:48:36.30
+13:30:45.80
+01:36:19.00
+01:16:38.20
+04:51:32.30
+10:55:31.10
+12:57:35.10
+11:47:11.40
+06:55:42.90
+08:57:51.20
+08:22:15.20
+10:27:16.90
+05:38:00.80
+07:03:16.60
+01:10:50.10
+05:29:29.03
+09:58:14.30
+03:45:34.80
+09:22:28.50
+05:56:43.90
+01:55:40.60
+08:19:26.60
+00:53:20.30
+26:21:45.60
+26:46:23.40
+24:42:30.10
+23:50:31.40
+23:58:29.30
+30:07:56.40
+25:24:46.80
+32:32:12.70
+06:42:20.10
+07:03:38.00
+08:56:20.10
+01:32:40.30
6
6
5
8
5
13
5
6
1
14
6
5
5
11
5
5
12
Ref.
72.34
− 130.23
− 512.71
− 154.6
− 67.84
86.22
− 123.05
− 213.24
281.48
137.3
− 185.13
− 369.47
− 248.94
57.71
− 174.75
93.61
379.37
− 609.61
37.29
− 110.69
− 78.89
− 83.02
− 85.78
− 125.96
− 14.9
− 205.16
− 57.78
− 76.28
− 151.56
− 41.6
− 122.92
− 146.63
− 84.43
− 5.54
− 30.06
129.07
347.86
128.88
242.18
47.32
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
17.15
14.58
12.41
13.16
12.12
12.43
12.48
17.99
19.56
15.42
12.32
15.11
10.8
13.11
12.06
13.5
10.5
12.68
13.24
16.45
15.18
12.05
12.84
11.49
12.06
12.8
11.17
10.79
12.61
17
11.46
10.05
12.4
14.78
13.66
9.63
16.86
41.56
13.57
21.43
μαcos δ
(mas yr−1 )
− 14.95
− 23.76
− 112.02
− 73.34
− 2.95
128.58
45.93
57.23
− 492.01
− 44.87
− 137.62
108.78
− 252.19
− 154.94
− 78.04
− 190.91
− 382.55
− 514.86
− 115.07
− 80.97
221.33
128.17
− 7.49
41.27
− 107.46
47.85
0.08
6.92
208.94
− 42.18
32.01
48.19
11.54
− 237.43
− 44.36
− 119.83
252.29
− 152.06
− 62.65
1.3
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
±
15.8
14.72
12.16
12.13
11.24
12.53
10.32
17.19
17.36
14.93
11.88
14.93
11.94
12.25
10.63
13.57
10.84
11.82
12.79
15.7
14.49
12.14
11.88
11.43
11.32
14
11.71
10.37
13.13
17.11
11.64
9.97
13.15
15.47
14.39
10.96
15.73
29.72
21.61
25.73
μδ
(mas yr−1 )
Dmin
(pc)
19.9
40.1
9.9
30.3
39.4
12.6
22
56.8
13.1
40.4
56.5
37.8
24.8
59.7
44.3
51.8
8.9
16.2
67.3
35.3
11.8
37.5
43.4
41.1
61.9
37.6
56
38.7
39.6
35.6
22.6
27.3
41.3
51.3
46.6
21.2
41.9
38.4
14.8
42.1
Baseline
(yr)
2.53
2.72
3.02
5.05
5.05
3.75
4.88
3.84
3.84
3.84
4.98
4.92
4.96
6.79
5.05
5.70
5.01
4.02
5.79
2.96
3.18
5.01
3.96
5.03
6.68
2.97
5.79
5.68
3.00
2.99
3.99
3.99
3.99
3.03
3.99
3.06
2.14W
0.78W
2.17W
1.28W
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
46.6
67.9
14.1
71.1
66.7
24.4
37.2
83.2
48.1
57.5
95.5
54.0
44.6
87.4
64.9
75.9
20.9
27.7
101.9
50.4
43.5
53.5
65.7
70.1
104.7
63.5
77.3
69.5
67
50.8
53.0
46.6
70.5
86.7
78.7
34
59.7
56.2
25.2
75.5
Dmax
(pc)
65.8
95.9
19.9
100.5
94.2
34.5
52.5
117.5
67.9
81.3
134.9
76.2
63.0
123.5
91.6
107.2
29.6
39.1
143.9
71.1
61.5
75.5
92.8
99.1
147.9
89.7
109.1
98.2
94.6
71.8
74.8
65.8
99.5
122.5
111.2
48.0
84.3
79.4
35.6
106.7
Dmax (binary)
(pc)
1: Chiu et al. (2006); 2: Knapp et al. (2004); 3: Kirkpatrick et al. (2011); 4: Burgasser et al. (2006); 5: Burningham et al. (2010b); 6: Pinfield et al. (2008); 7: Lodieu et al. (2007b); 8: Albert et al. (2011);
9: Burningham et al. (2008); 10: Burningham et al. (2011); 11: Tsvetanov et al. (2000); 12: Burningham et al. (2010a); 13: Kendall et al. (2007); 14: Murray et al. (2011).
α
(J2000)
Name
Table 5 – continued
76 T dwarfs
477
478
B. Burningham et al.
Although estimating M dwarf metallicities is challenging, the
good agreement of the different estimates that we have obtained for
LHS 6176A highlights the excellent progress that has been made
in this field in recent years. We adopt the mean of the estimates
from Neves et al. (2012), Mann et al. (2013) and Rojas-Ayala et al.
(2012), and assign a metallicity of [Fe/H] = −0.3 ± 0.1 dex for
LHS 6176A. Table 7 summarizes our determined properties for
LHS 6176A.
6.2.3 The properties of LHS 6176B
LHS 6176A. Johnson V-band data were obtained using the 50 cm
pt5m telescope on La Palma on the night of 2012 December 17.
56 s exposures were obtained. These exposures had the dark current
and bias levels subtracted and were flat fielded using twilight sky
frames. Objects were detected and instrumental magnitudes calculated using SEXTRACTOR (Bertin & Arnouts 1996). The instrumental
magnitudes were calibrated against V-band magnitudes from the
AAVSO Photometric All-Sky Survey6 . The resulting transformation was V = 0.9774 × I + 0.1010, where V is the calibrated
Johnson V-band magnitude and I is the instrumental magnitude.
The photon noise in our measurement is 0.004 mag, but this is outweighed by approximately 0.03 mag of calibration error, which we
instead quote as our uncertainty (see Table 7).
We thus estimate a moderately low metallicity of
[Fe/H] = −0.43 ± 0.19 for LHS 6176A from the Schlaufman
& Laughlin (2010) calibration and [Fe/H] = −0.36 ± 0.17 dex
using the Neves et al. (2012) calibration. The quoted error reflects
the dispersion about the metallicity relations which dominates over
our photometric uncertainties.
Method 2. We have applied the method of Mann et al. (2013) to
our SNIFS optical and our JHK SpeX NIR spectra. From the optical
regions we estimate [Fe/H] = −0.31 ± 0.08 dex, whilst the NIR
regions yield an estimate of [Fe/H] = −0.25 ± 0.05 dex. These
are consistent with one another, and the optically based estimate is
consistent with the photometric estimate based on the Neves et al.
(2012) relations.
Method 3. Using the strengths of metal sensitive K-band features
and the calibration described by Rojas-Ayala et al. (2012) provides
an estimate of [Fe/H] = −0.26 ± 0.14 dex. This is consistent
with all the other estimates, and matches the other spectroscopic
estimates particularly well.
6
http://www.aavso.org/apass
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 7. A V − J versus MJ colour magnitude diagram showing the Hipparcos stars with V and J photometry (black dots) and our candidate primary
stars. Blue symbols indicate the hypothetical range of MJ for primaries selected from LSPM, whilst the red symbols indicate the same for candidate
primaries selected from NOMAD. One candidate primary lies beyond the
plotted range.
We have used our derived properties for LHS 6176A to constrain
the metallicity and age of LHS 6176B, and determine more precise properties than would otherwise have been possible. Since our
warm-Spitzer photometry is more precise than the WISE survey
photometry, we have used it in preference in our calculations. We
followed the same method as that described in Burningham et al.
(2011) and Pinfield et al. (2012).
Briefly, we have calculated the bolometric flux using our YJHK
spectroscopy flux calibrated to our J-band photometry and warmSpitzer photometry. We have filled the gaps in our wavelength coverage by scaling the latest BT Settl model spectra (Allard, Homeier
& Freytag 2010) to match our Y-band spectrum for wavelengths
blueward of our GNIRS spectrum, and to match our Ch1 and Ch2
photometry for wavelengths redward of it. To avoid biasing our
derived flux by assumptions regarding the Teff of the target, we
initially calculated the flux using a wide range of models covering
Teff = 500–1000 K, log g = 4.5–5.5 and metallicity, [M/H] = 0.0
and −0.3. The resulting flux ruled out high- and low-Teff extremes
and lowest gravity cases. We thus recalculated the flux using Teff =
600–900 K models with log g = 5.0–5.5.
To account for random uncertainties in our flux calibration, we
calculated each flux estimate (for each model spectrum) as the
mean of a set of 100 different scalings, each offset by a random
value drawn from the uncertainty in the photometry. Our final flux
estimate is the mean of the estimates made using the range of
models and different scalings, and our error is taken as the standard
deviation on this value. Thus, our uncertainty implicitly includes
both random and (identified) systematic elements. We thus calculate
the flux from LHS 6176B as Fbol = 2.13 ± 0.15 × 10−16 W m−2 ,
and its luminosity as 8.94 ± 1.30 × 1020 W or log(L∗ /L ) =
−5.63 ± 0.07. This is approximately 60 per cent higher than the
luminosity of the similarly metal poor benchmark T8 dwarf BD+01
2920B which was calculated using the same method (Pinfield et al.
2012).
We can use the measured luminosity for LHS 6176B, in combination with the evolutionary models of Baraffe et al. (2003), to
estimate its radius and mass, assuming our estimated age for LHS
6176A. The COND models assume solar metallicity, and the effect of low metallicity on the luminosity–radius relation of brown
dwarfs is not well constrained by observations. However, theoretical correlations between metallicity and radius derived by Burrows,
Heng & Nampaisarn (2011) suggest radii are reduced by less than
5 per cent for a −0.5 dex shift from solar metallicity for objects
at the stellar/substellar boundary, with considerably smaller shifts
at lower masses. The Saumon & Marley (2008) evolutionary sequences suggest that shifts in radius of less than 1 per cent can be
expected as a result of decreasing metallicity by 0.3 dex.
The derived properties also depend strongly on the assumed
multiplicity (or otherwise) of LHS 6176B. If we assume that
LHS 6176B is a single object, with an age in excess of 3.5 Gyr
(based on the lack of Hα emission seen in LHS 6176A), we
76 T dwarfs
479
Table 6. Initial characterization of candidate wide binary pairs selected by our cross-matches against the LSPM, NOMAD and Hipparcos
catalogues. The ‘Notes column includes previous discovery references and chance alignment probabilities for new candidates that passed the
common distance test.
T dwarf
Primary name
Separation (arcsec)
μαcos δ
μδ
V−J
Distance match?a
Notes
ULAS J0950+0117
ULAS J1043+1048
ULAS J1300+1221
ULAS J1315+0826
ULAS J1335+1130
ULAS J1339+0104
ULAS J1423+0116
ULAS J1459+0857
ULAS J1504+0538
LHS 6176
NOMAD 1006−0190624
Ross 458
NOMAD 0983−0263649
LSPM J1334+1123
HD 118865
BD+01 2920
LSPM J1459+0851
HIP 73786
52
1021
105
382
1625
148
153
386
64
0.5σ
3.9σ
0.3σ
1.6σ
1.0σ
2.3σ
2.9σ
0.6σ
0.1σ
3.5σ
0.7σ
0.2σ
2.0σ
2.5σ
1.6σ
0.8σ
1.6σ
0.8σ
4.11
1.2
3.28
1.3
3.11
1.0
1.21
–
2.64
Y
N
Y
B
N
Y
Y
–
Y
0.0002 per cent
1
2 per cent
0.3 per cent
2
3
4,5
= good match; N = bad match; B = requires binarity.
(1) Goldman et al. (2010); (2) Pinfield et al. (2012); (3) Day-Jones et al. (2011); (4) Scholz (2010); (5) Murray et al. (2011).
aY
LHS 6176A
RA (ep = 2011.2579 eq = 2000)
Dec. (ep = 2011.2579 eq = 2000)
PMαcos δ (mas yr−1 )
PMδ (mas yr−1 )
Spectral type
V
BJ
J (2MASS)
J − H (2MASS)
H − Ks (2MASS)
V − Ks
J (UKIDSS)
Y − J (UKIDSS)
J − H (UKIDSS)
H − K (UKIDSS)
W1
W2
W3
W4
[3.6]
[4.5]
π
Distance
[Fe/H]
Hα EW
Age
log(L∗ /L )
Projected separation
LHS 6176B
09:50:49.8
09:50:47.3
+01:18:09.4
+01:17:33.0
242.42 ± 19.0
237.18 ± 2.84
−351.46 ± 4.50
−360.03 ± 3.13
M4
T8p
13.88 ± 0.03
–
15.2a
–
9.80 ± 0.02b
–
0.57 ± 0.04b
–
0.28 ± 0.04b
–
4.93 ± 0.04
–
–
18.02 ± 0.03
–
0.88 ± 0.04
–
−0.38 ± 0.04
–
−0.45 ± 0.08
8.77 ± 0.02
18.05 ± 0.34
8.60 ± 0.02
14.48 ± 0.06
8.50 ± 0.02
>12.85
>7.98
>9.20
–
16.28 ± 0.01
–
14.35 ± 0.02
46.14 ± 10.7
53.40 ± 3.51
21.7+6.5
pc
18.731.32
−4.1
−1.15
−0.30 ± 0.1
–
−0.29 ± 0.23 Å
–
>3.5 Gyrc
–
–
−5.63 ± 0.07
52 arcsec, ∼970 au
Figure 8. Our SNIFS spectrum of LHS 6176A compared with the nonactive M4 template spectrum from Bochanski et al. (2007).
a Lépine
& Shara (2005).
2MASS All-Sky Point Source Catalog.
c Derived from activity lifetime information presented in West et al. (2008).
b From
estimate its radius to be 0.078 R < R < 0.094 R mass to be
0.055 M > M > 0.030 M , with a corresponding temperature of
850 > Teff > 710 K and gravity 5.3 > log g > 5.0 (respectively).
If, on the other hand, we assume that LHS 6176B is an equal
luminosity binary system, we find that the components would have
radius 0.081 R < R < 0.096 R and mass 0.045 M > M >
0.022 M , and with a corresponding temperature of 700 > Teff >
590 K and gravity 5.30 > log g > 4.8. Although LHS 6176B has a Jband magnitude (MJ = 16.65) at the faint end of the scatter about the
T8 mean magnitude in Dupuy & Liu (2012, MJ = 16.39 ± 0.35),
we cannot rule out binarity. It should be borne in mind that the
mean T8 MJ is calculated from a sample that specifically excludes
Figure 9. Our JHK SpeX spectrum of LHS 6176A.
peculiar objects, and so is dominated by objects with higher metallicity than LHS 6176B. A comparison with other low-metallicity
objects would thus be more relevant. BD+01 2920B (Pinfield et al.
2012) provides just such a comparison. LHS 6176B is approximately 0.85 mag brighter than this object the J band. It is thus quite
possible that LHS 6176B is an unresolved binary system, and we
refrain from adopting a single set of properties for this object.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Table 7. Properties of LHS 6176AB.
480
B. Burningham et al.
Table 8. Properties of HD 118865AB.
HD 118865A
13 39 34.33a
13:39:43.79
+01 05 18.12a
+01:04:36.40
−95.75 ± 0.80a
−130.23 ± 14.58
−48.81 ± 0.54a
−23.76 ± 14.72
F5
T5
8.52 ± 0.02b
–
7.98 ± 0.01b
–
–
6.98 ± 0.02c
0.25 ± 0.05c
–
0.06 ± 0.05c
–
–
18.08 ± 0.04
–
1.07 ± 0.06
–
−0.31 ± 0.14
–
0.0 ± 0.13
–
16.93 ± 0.01
–
16.08 ± 0.02
16.02 ± 0.86a
–
62.4+3.2
–
−3.6 pc
0.09 ± 0.09d
–
1.5−4.9 Gyr (1σ interval)d
–
−5.24 ± 0.04
148 arcsec, ∼9200 au
a van
Leeuwen (2007); b Høg et al. (2000); c From 2MASS AllSky Point Source Catalog; d Casagrande et al. (2011).
6.3 HD 118865 AB
The probability of chance alignment for HD 118865A and ULAS
J1339+0104 is 0.2 per cent, which is significantly higher than has
been found previously for wide binary systems in our searches
(e.g. Burningham et al. 2009; Pinfield et al. 2012). This is due to
the combination of a large range of plausible distances for the T5
dwarf (40–95 pc, allowing for the possibility that it is a binary)
and the relatively low proper motions of the proposed components.
However, if our range of plausible distances is reduced to account
only for the spread in T5 absolute magnitudes (rather than T4.5–
T5.5 and binarity; since the former arguably already incorporates
these effects), then we find the probability of chance alignment is
reduced to 0.001 per cent. On balance, it is reasonable to proceed
with analysis of this system as a bona fide common proper motion
binary system, although we caution that a parallax for the T dwarf
is required to confirm beyond doubt that this pair is associated.
In many respects, HD 118865A represents an ideal benchmark
primary star. With a Hipparcos parallax (Perryman et al. 1997;
van Leeuwen 2007), kinematic and model-based age estimates and
well-measured metallicity (Casagrande et al. 2011), it avoids many
of the pitfalls associated with less massive primary stars. Indeed,
the metallicity diagnostics for the M dwarf primaries that dominate
the substellar wide binary sample have been benchmarked against
nearby binary systems containing FGK primary stars for these reasons. The properties of HD 118865A are summarized in Table 8.
To determine the properties of HD 118865B, we have followed
an identical method to that applied for LHS 6176B (see Section
6.2.3). In this case, we have used models spanning the Teff = 1000–
1400 K range, log g = 4.5–5.5 and solar metallicity. After a first
round of calculations, we found that the lowest gravity models were
inconsistent with our estimated luminosity and age for the system,
and so were excluded from the next iteration. We find the flux from
HD 118865B to be Fbol = 1.47 ± 0.14 × 10−16 W m−2 and the
luminosity to be Lbol = 6.85 ± 1.00 × 1021 W or log(L∗ /L ) =
−4.75 ± 0.07.
6.4 HIP 73786B
Scholz (2010) and Murray et al. (2011) independently identified
HIP 73786AB as a wide binary system consisting of a metal poor
K5 dwarf moving in common motion with a T6p dwarf, and we
recovered it with our search. At the time of its discovery longer
wavelength photometry for the T dwarf was not available, and so
no bolometric flux estimate was made. The advent of the WISE
all-sky release, however, allows such an estimate to made relatively
conveniently and for the purposes of including this object in subsequent discussion we have estimated the properties for HIP 73786B
following the same method as described in the previous sections.
In this case, we have used BT Settl models spanning the Teff =
800–1200 K range, log g = 4.5–5.5 and solar metallicity, scaling
the longer wavelength portions to match the W1 and W2 survey
photometry, and stitching them to our flux-calibrated NIR spectrum
at 2.4 µm. As before, the regions blueward of 1 µm have been
filled using models scaled to match the Y-band flux in our NIR
spectrum. The current lack of low-metallicity BT Settl models in
this temperature range has prevented us from using models covering
the entire expected parameter range for this object, but since we
scale the models to match the long-wavelength photometry, we do
not expect this to have a significant impact on our flux estimate.
We find the flux from HIP 73786B to be Fbol = 5.31 ± 0.30 ×
10−16 W m−2 and the luminosity to be Lbol = 2.20 ± 0.12 × 1021 W
or log(L∗ /L ) = −5.24 ± 0.04.
Murray et al. (2011) loosely constrained the age of the HIP
73786AB system to be 1.6–10 Gyrs. If we accept this age for
HIP 73786B, then from its luminosity and the COND evolutionary models we infer a radius of 0.076 R < R < 0.096 R and a
mass of 0.063 M > M > 0.028 M , with corresponding temperature 1020 > Teff > 910 K and gravity 5.5 > log g > 4.9 under the
assumption that it is a single object. If HIP 73786B is an equal luminosity binary then we derive a radius 0.076 R < R < 0.098 R ,
0.058 M > M > 0.022 M , 860 K > Teff > 760 K and 5.4 >
log g > 4.8.
6.5 Properties of the UKIDSS wide binary sample
The relatively uniform manner in which our wide binary systems
have been identified allows us to draw some preliminary conclusions
about the properties of the late-T dwarf wide binary companion population. First, it is apparent that with 7 wide binary companions out
of a total of 92 T dwarfs within our proper motion >100 mas yr−1
selection, we can place a minimum value of 8 per cent on the wide
binary companion fraction. This is consistent with the minimum
value of 5 per cent found by Gomes et al. (2013) for L dwarfs.
This represents a lower limit, as our selection of candidate primary
stars is limited by available photometry for faint red primary star
candidates, as highlighted in Section 6.1. This source of incompleteness also explains another feature of our binary sample. That
is, that the latest type T dwarfs are significantly over-represented,
with 25 per cent of T8 and later objects (3 out of 12 in our selection)
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RA (J2000)
Dec. (J2000)
PMαcos δ (mas yr−1 )
PMδ (mas yr−1 )
Spectral type
BT
VT
J (2MASS)
J − H (2MASS)
H − Ks (2MASS)
J (UKIDSS)
Y − J (UKIDSS)
J − H (UKIDSS)
H − K (UKIDSS)
[3.6]
[4.5]
π
Distance
[Fe/H]
Age
log(L∗ /L )
Projected separation
HD 118865B
If we assume that HD 118865B is a single object with an age
of between 1.5 and 4.9 Gyr, then we find that its luminosity corresponds to a radius of 0.080 R < R < 0.091 R and a mass
of 0.065 M > M > 0.040 M , with corresponding temperature
1320 K > Teff > 1240 K and gravity 5.4 > log g > 5.0, according
to the COND evolutionary models. If HD 118865B is an equalmass binary, these properties should be revised to 0.079 R < R <
0.095 R , 0.060 M > M > 0.030 M , 1120 K > Teff > 1020 K
and 5.4 > log g > 5.0.
76 T dwarfs
6.6 The colours of benchmark T dwarfs and the latest model
atmospheres
Fig. 10 shows H − K versus H − [4.5] colour–colour plots for
the compendium of MKO and Spitzer photometry of late-T dwarf
benchmark systems including that presented by Leggett et al.
(2010a) and updated with additional photometry presented here.
Fig. 11 shows similar Y − J versus J − W2 colour–colour diagrams
for benchmark T dwarfs with WISE photometry. The properties of
the benchmarks are summarized in Table 9.
Overlaid on the two panels of Figs 10 and 11 are colour tracks
for the models of Saumon et al. (2012) and Morley et al. (2012),
and a 1σ box for the model-predicted colours of each benchmark
is connected to each objects observed colours (where a model prediction is available on these grids). These models incorporate the
latest NH3 opacities from Yurchenko, Barber & Tennyson (2011),
and a new treatment of collisionally induced H2 absorption (CIA
H2 ; Abel et al. 2011, 2012). The Morley et al. (2012) models also
include the effects of proposed sulphide and alkali condensates,
which may become important in atmospheres below Teff ≈ 800 K.
Low-metallicity versions of these models are not yet available, so it
is not possible to assess if these recent developments have impacted
the predictions for low-metallicity atmospheres.
It is difficult to draw firm conclusions about the success or otherwise of the different models by comparing the colour predictions to
the observed colours of the benchmark systems as deficiencies in the
models can be both masked and amplified by their differing impacts
at the wavelengths of the two photometric bands being compared
in any particular case. Nonetheless some of the principal differences between the models and the observations can be attributed to
known deficiencies in the model grids. For example, the predicted
H − [4.5] and J − W2 colours for the solar-metallicity benchmarks
are almost universally too red in both Figs 10 and 11. This is due to
the fact that neither model set includes non-equilibrium chemistry,
which has been demonstrated to be important for determining the
emergent flux in the 4−5 µm region (e.g. Saumon et al. 2006).
The most obvious conclusion that can be drawn from examining
the colours of the benchmark systems is that metallicity and gravity
both have a significant impact on the near to mid-infrared colours
of late-T dwarfs, along with the (generally considered) dominant
influence of temperature. For example, the low-gravity benchmark
Ross 458C appears only 0.15–0.2 mag redder in H − [4.5] than Gl
570D and HD 3651B, despite being approximately 100K cooler.
Comparing the colours of Wolf 940B with Gl 570D and HD 3651B
suggests that Teff = −100 K should (approximately) correspond
to (H − [4.5]) = +0.6 mag. This implies that the low gravity of
Ross 458C imposes a blueward shift of ∼0.4 mag for log g ≈
−0.7 dex, or roughly +0.1 mag per +0.18 dex in log g.
The impact of metallicity appears to be even more significant,
with lower metallicity objects appearing redder (in e.g. H − [4.5]
or J − W2) at a given Teff than objects with near-solar metallicity.
For example, comparing the colours of LHS 6176B and BD+01
2920B with the warmer solar metallicity objects suggest that a shift
of [Fe/H] ≈ −0.3 dex can give rise to a shift of (H − [4.5]) ≈
+0.6, if we assume that LHS 6176B is an unresolved binary with
a Teff = 645 ± 55 K and log g = 5.05 ± 0.25. If LHS 6176B is a
single object then the apparent shift in colour due to its metallicity
would more like 0.9 mag. This is somewhat higher than for BD+01
2920B, and further argues for a binary interpretation for this object.
This effect has been suggested previously, based on comparisons
of photometric colours and atmospheric models (Leggett et al. 2009,
2010a), and can largely be attributed to increased flux in the 4.5 µm
region due to reduced CO opacity. Comparison of solar and lowmetallicity cases for the BT Settl (Allard et al. 2010) models that
were first shown in Pinfield et al. (2012) and the Saumon & Marley (2008) models indicates that although both predict a colour
shift due to lower metallicity, both underpredict its magnitude. The
Saumon & Marley (2008) models predict (H − [4.5]) ≈ +0.2
for [M/H] = −0.3, whilst the BT Settl grid predicts
(H − [4.5]) ≈ +0.3 for the same change at Teff = 700 K, compared
with the (H − [4.5]) ≈ +0.6 shift seen in our benchmarks.
The slightly stronger metallicity dependence that the BT Settl
model colours exhibit can likely be attributed to the combination
of two factors. First, the BT Settl models include non-equilibrium
chemistry for CO2 , which will result in a greater relative increase
in flux in the [4.5] band in response to reducing the metallicity.
Secondly, the BT Settl models include additional methane opacity
in the H band, where the methane line lists are very incomplete,
based on a statistical estimate of the contributions from the hot
vibrational bands. With decreased metallicity these hot bands are
likely to become more important as the chemistry shifts in favour
of CH4 over CO and CO2 (e.g. Lodders & Fegley 2002). Indeed,
it is likely that the failure of the models to accurately reproduce
the strong metallicity dependence of the H − [4.5] colour can be
partially attributed to the incomplete nature of the methane line lists.
The strong dependence on metallicity of the H − [4.5] and J − W2
colours should be considered carefully when using these colours to
estimate the properties of cool brown dwarfs. Two specific examples
worth highlighting in this context are: SDSS 1416+1338B, the
second reddest known T dwarf in H − [4.5] despite a spectral type
of T7.5 (Burningham et al. 2010a); and WISEPC J1828+2650,
which is the reddest known Y dwarf in the same colours (Cushing
et al. 2011; Leggett et al. 2013). In the case of the former, its
anomalously red colour was initially interpreted as being indicative
of Teff ≈ 500 K (Burningham et al. 2010a). However, its parallax
has since been measured, and its luminosity appears to rule out
such a low temperature, and its extreme colours are now attributed
to some combination of low metallicity and/or high gravity with
a significantly higher Teff (Murray et al., in preparation). In the
case of WISEPC J1828+2650, Leggett et al. (2013) argues that its
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appearing as wide binary companions, compared to 5 per cent for
T4.5–T6.5 dwarfs. Whether this suggests that the true wide binary
companion fraction for T dwarfs is nearer the 25 per cent value
seen for the latest type objects is an open issue that will require
significantly improved selection of earlier type binary systems.
The scarcity of earlier type wide binary companions can be understood not in terms of bias within our proper motion sample,
but rather as a reflection of the relative space densities of late- and
mid-type T dwarfs, and the bias against fainter red candidate primary stars with Tycho photometry. The space density of the T8
and later dwarfs is significantly (e.g. four times) higher than that
of mid-type T dwarfs. This means that a larger proportion of the
latest type objects will be found at close distances, where the LSPM
and NOMAD catalogues are relatively complete for primary stars
and proper motions are larger. The earlier type objects, however,
will be less numerous in the nearby volume, and so larger fraction
will be found at larger distances, where the catalogues of potential
primary stars are most incomplete for the most common M dwarf
type of primaries, and proper motions are smaller making reliable
identification more problematic. It thus appears, at this stage, that
we can (weakly) conclude that the spectral type distribution of T
dwarfs as wide binary companions does not appear to be drastically
different from that of isolated (i.e. single) T dwarfs.
481
482
B. Burningham et al.
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Figure 10. H − [4.5] versus H − K colour–colour plots for the compilation of photometry for T dwarfs from Leggett et al. (2010a) and benchmark systems
from the literature, along with new photometry and benchmarks presented in this paper. Benchmark systems with roughly solar metallicity are shown in green,
whilst those with subsolar metallicity are shown in blue. Model colour tracks for Saumon et al. (2012), top panel, and Morley et al. (2012), bottom panel,
are shown for comparison. The latter assumed fsed = 5. The models shown are all solar metallicity, and the Teff and log g values are indicated on the colour
sequences. Each benchmark object is linked to a 1σ box indicating the model prediction for its colours.
76 T dwarfs
483
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Figure 11. J − W2 versus Y − J colour–colour plots for the UKIDSS T dwarfs with YJ and WISE photometry and benchmark systems from the literature, along
with new photometry and benchmarks presented in this paper. Benchmark systems with roughly solar metallicity are shown in green, whilst those with subsolar
metallicity are shown in blue. Model colour tracks for Saumon et al. (2012), top panel, and Morley et al. (2012), bottom panel, are shown for comparison. The
latter assumed fsed = 5. The models shown are all solar metallicity, and the Teff and log g values are indicated on the colour sequences. Each benchmark object
is linked to a 1σ box indicating the model prediction for its colours.
484
B. Burningham et al.
Table 9. Summary of the properties of the benchmark systems shown in
Figs 10 and 11.
Name
Wolf 940B
BD+01 2920B
Ross 458C
LHS 6176B
Gl 570D
HD 3651B
HIP 73786B
LHS 2803B
HD 118865B
Teff (K)
605
680
695
780
800
810
965
1120
1280
±
±
±
±
±
±
±
±
±
20a
55c
60d
70e
20f
30g
55e
80i
40e
5.0
5.0
4.35
5.15
5.1
5.3
5.2
5.4
5.2
log g
[Fe/H]
±
±
±
±
±
±
±
±
±
+0.02 ± 0.05b
− 0.36 ± 0.06c
+0.09 ± 0.05b
− 0.30 ± 0.10e
+0.09 ± 0.04f
+0.12 ± 0.04g
− 0.30 ± 0.1h
∼0.0i
0.09 ± 0.10e
0.2a
0.3c
0.35d
0.15e
0.15f
0.2g
0.3e
0.1i
0.2e
c Pinfield
extremely red H − [4.5] can only be consistent with its luminosity
(which is higher than for some earlier and bluer Y dwarfs) if it is
either younger than 50 Myr (with a mass <1MJup ) or an unresolved
binary system. Given the example of SDSS J1416+1348B and the
colours of benchmark systems in Fig. 10 (albeit at higher Teff ),
it seems reasonable to also offer a third (somewhat speculative)
interpretation e.g. this object may in fact be somewhat warmer and
more massive, but with its colours reddened by low metallicity, high
gravity and/or their combined impact on cloud properties.
It is important to emphasize that what follows is based on a
highly simplistic extrapolation of the trends we have identified at
cool T dwarf temperatures (Teff = 600−800K) to Y dwarf temperatures some 200–500 K cooler. We note that at these low temperatures the gravity range is limited by the long cooling times for
the most massive objects (e.g. Saumon & Marley 2008), so the
7 T H E D I S T R I B U T I O N O F T DWA R F
COLOURS
Beyond comparisons to our benchmark sample, it is also interesting to make more qualitative comparisons between the spread in
colours predicted by the atmospheric models and that seen for the
wider T dwarf population. Such comparisons provide insight as to
whether the impact of varying parameters such as gravity in models
results in colour shifts of the similar proportions seen in the data.
Figure 12. H − [4.5] versus J − H colour–colour plots for the compilation of photometry for T dwarfs from Leggett et al. (2010a), along with new photometry
presented in this paper. Benchmark systems are indicated and numbered as for Fig. 10.
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et al. (2010b); b From the V − K calibration of Neves et al. (2012);
et al. (2012); d Burningham et al. (2011); e This work; f Saumon
et al. (2006); g Liu, Leggett & Chiu (2007); h Murray et al. (2011); i Deacon
et al. (2012).
a Leggett
bulk of any colour shift at a given Teff must be driven by metallicity. If we apply the shift that we have identified at higher Teff of
H − [4.5] = +0.6 mag for roughly [M/H] = −0.3 dex, and
assume that the shift is entirely driven by extra flux in [4.5], we
find that a halo-like low metallicity of [M/H] ≈ −1 could provide
the additional 1.5−2.0 mag of flux seen in the 4.5 µm region from
WISEPC J1828+2650 (Beichman et al. 2013; Leggett et al. 2013).
Although the tangential velocity of WISEPC J1828+2650 has been
measured as 51 ± 5 km s−1 (Beichman et al. 2013), which is most
consistent with thin-disc membership, this does not rule out a metallicity as low as [M/H] ≈ −1. Its NIR colours may also lend weight
to the low-metallicity interpretation. Leggett et al. (2013) report
H − K = −0.63 ± 0.43 for this target, making it the bluest
known Y dwarf in these colours, which at higher temperatures
would be associated with low metallicity and/or high gravity. Its red
J − H = +0.63 ± 0.33 colour is also not at odds with low metallicity. For example, in the T dwarf regime metallicity has little impact
on the J − H colours of T dwarfs (see Fig. 12). Interestingly, the
BT Settl models predict redder J − H colours for lower metallicity
T dwarfs, with a more pronounced effect at lower Teff . It will be interesting to see how the picture of the impact of varying metallicity
develops as new models grids are calculated and new benchmark
systems are discovered over the coming years.
76 T dwarfs
8 K I N E M AT I C S O F T H E U K I D S S L AT E - T
DWA R F S A M P L E
The implication from Figs 10–14 that the late-T dwarf sample is
not dominated by young low-mass objects as had been previously
suggested (Leggett et al. 2010a; Burningham et al. 2011; Pinfield
et al. 2012) highlights the inherent model dependency of deter-
mining the properties of cool brown dwarfs from colour–colour
diagrams. To shed further light on this issue, we have constructed a
J-band-reduced proper motion diagram for our sample (see Fig. 15).
It is apparent that there are no obvious features which distinguish
the distribution of kinematic properties between the latest type T
dwarfs and the earlier type objects in UKIDSS. A more robust statistical treatment to assess this will be presented by Smith et al. (in
preparation). Nonetheless, it appears that the late-T dwarf sample is
kinematically indistinct from the mid-T dwarfs, which themselves
have been shown to match the kinematics of the earlier type L and M
dwarfs, reflecting the typical Galactic disc age distribution (Faherty
et al. 2009).
We note that our earliest type objects, in the T4 class, appear
to be preferentially distributed to higher speeds. Previous detailed
studies of T dwarf kinematics (e.g. Faherty et al. 2009) argue this
is not a real effect. It likely arises as a result of the combination of
our J − H < 0.1 and J − K < 0.1 colour requirements, which may
exclude T4 dwarfs with redder H − K colours, and J − H colours
near our colour cut-offs. As such, we are likely to preferentially
select objects that are bluer in H − K, due to high gravity and/or
low metallicity, in this region. Such objects can be expected to be
typically older, and thus will exhibit higher velocities.
9 U P DAT E D S PAC E D E N S I T Y E S T I M AT E
We now have near-complete follow-up of all candidates with
J < 18.8 in UKIDSS LAS DR8. This covers 2270 square degrees
of sky within the SDSS DR8 footprint. Our significantly increased
sample of ≥T6 dwarfs allows us to improve on the space density
estimate derived in Burningham et al. (2010b). However, the optimization of our selection method since that work (to include CH4
imaging) somewhat complicates our completeness correction and
bias selection. Unifying a sample selected in a slightly inhomogeneous manner, such as ours, to derive a space density is best
achieved via a Bayesian parameter estimate method. For this reason, we defer this more rigorous derivation of the space density to
a future work, since it is beyond the scope of this discovery paper.
Here, instead, we follow the method of Pinfield et al. (2008) and
Burningham et al. (2010b), and derive approximate correction factors for the incompleteness introduced by our different photometric
cuts in the same manner as we previously did for correcting for our
J − H selection cut, and simply scale these by the proportion of the
sample to which they were applied.
There are 76 dwarfs with measured spectral types of T6 or later
with J < 18.8 in the region of sky covered by UKIDSS DR8 and
SDSS DR8 probed by our searches. Of these, 39 have already had
spectra published (Tsvetanov et al. 2000; Burgasser et al. 2004;
Chiu et al. 2006, 2008; Lodieu et al. 2007b; Pinfield et al. 2008;
Burningham et al. 2008, 2009, 2010b; Delorme et al. 2008; Kirkpatrick et al. 2011; Scholz et al. 2012), and 37 have been presented for the first time here. In addition, we have four targets with
CH4 s − CH4 l < −0.5 which have not yet been followed up spectroscopically (see Table 1). The CH4 types for two of these objects
are constrained to earlier than T5.5, whilst two are consistent with
types as late as T6. We thus include the latter two of these four in
our sample, bringing it to a total of 78 T6 and later dwarfs.
We also have one target for which we do not have wellcalibrated methane photometry and no spectroscopic observations: ULAS J 0800+1908 has a LIRIS methane colour of
H − CH4 l = −0.73. This is suggestive of a type intermediate
between that of the T7 dwarf ULAS J0746+2355 which has LIRIS
H − CH4 l = −0.87 and the T6.5 dwarf ULAS J1023+0447.
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Fig. 10 shows H − K versus H − [4.5] colour–colour plots for the
compendium of MKO and Spitzer photometry of late-T dwarfs presented by Leggett et al. (2010a), updated with additional photometry
presented here. The benchmark T dwarfs are labelled as for Fig. 10.
Fig. 11 shows similar Y − J versus J − W2 colour–colour diagrams
for the UKIDSS T dwarfs with WISE photometry. These plots reveal that the two model sets are each able to reproduce the spread
in late-T dwarf colours in one of H − K or Y − J, but not both. We
note that neither set of models include refractory condensate clouds
that are important at warmer Teff , and so they are not expected
to match the data for earlier type T dwarfs, which they do not.
The upper panel of Fig. 13 suggests that the revised CIA
H2 opacity included in the Saumon et al. (2012) models has
significantly improved the match to the spread and pattern seen
in H − K colours over previous model grids, which tended to predict H − K that was too blue, and descending to more negative
values rapidly with Teff (e.g. Burningham et al. 2009; Pinfield et al.
2012). The lower panel of Fig. 13 shows the colour tracks for the
Morley et al. (2012) models, which, in addition to using the Saumon
et al. (2012) CIA H2 opacity also include alkali and sulphide clouds
that may become important at Teff 800 K. These models show
slightly redder H − K, with a smaller spread in H − K colour as a
function of gravity. The reduced spread in H − K arises as higher
gravity leads to thicker cloud layers, pushing the colours redward
and partially counteracting the blueward trend that would otherwise
result from the increased CIA H2 opacity.
As can be seen in Fig. 14, the introduction of the low-Teff clouds
tends to increase the Y − J spread due to gravity, as the effect
of increasing gravity on pressure sensitive gas-phase opacities and
the impact of thickening cloud layers with increasing gravity both
tend to give redder Y − J colours, and here the cloudy models better
reproduce the observed spread in colours than those without clouds.
Variations in metallicity may have a significant impact on cloud
properties. Since only solar metallicity realizations are available
for the new model grids, it is thus impossible to determine whether
their failure to reproduce the observed colour spread in both Figs 13
and 14 merely reflects this fact, or if it suggests some other as-yet
unidentified shortcoming in the theoretical approach.
We have presented plots similar to Fig. 13 in previous work (e.g.
Leggett et al. 2010a; Burningham et al. 2011; Pinfield et al. 2012)
comparing the colours of the late-T dwarf sample with the previous
generation of models that did not include the improved CIA H2
treatment, nor the low-temperature sulphide and alkali condensates
proposed by Morley et al. (2012). The previous generations of
models typically predicted significantly bluer colours in H − K,
with H − K becoming increasingly blue with decreasing Teff . The
preference for redder H − K colours amongst the late-T sample
led us to conclude that there was some bias in favour of young,
low-gravity objects either within our selection or in the population
itself. The model colour tracks in the top panel of Fig. 13 match the
observed spread in H − K colours very well, with no need for such a
young-age bias in the sample. Indeed, when both Figs 13 and 14 are
considered, the strong model dependence of previous conclusions
drawn from such plots is apparent.
485
486
B. Burningham et al.
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Figure 13. H − [4.5] versus H − K colour–colour plots for the compilation of photometry for T dwarfs from Leggett et al. (2010a), along with new photometry
presented in this paper. Model colour tracks for Saumon et al. (2012), top panel, and Morley et al. (2012), bottom panel, are shown for comparison, the latter
assumed fsed = 5. The models shown are all solar metallicity, and the Teff and log g values are indicated on the colour sequences. Benchmark systems are
indicated and numbered as for Fig. 10.
76 T dwarfs
487
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Figure 14. J − W2 versus Y − J colour–colour plots for the UKIDSS T dwarfs with YJ and WISE photometry. Model colour tracks for Saumon et al. (2012),
top panel, and Morley et al. (2012), bottom panel, are shown for comparison, the latter assumed fsed = 5. The models shown are all solar metallicity, and the
Teff and log g values are indicated on the colour sequences. Benchmark systems are indicated and numbered as for Fig. 11.
488
B. Burningham et al.
However, the lack of a spectral type calibration for this methane
colour prevents us from assigning this object to a specific spectral
type bin in our sample. Instead, we incorporate it into the uncertainty
in the T6 and T7 bins.
Two-thirds of our sample were selected using the method presented in Pinfield et al. (2008) and Burningham et al. (2010b). For
this method, we found that a correction factor of 1.03 should be
applied to all spectral type bins, to account for the biases introduced
by the following effects: scatter out of our J − H < 0.1 selection
due to photometric error; mismatching of bona fide T dwarfs with
background stars in our SDSS cross-match; wide binary companions to stars (which should be excluded from a T dwarf primary
space density). We do not correct for the possibility of excluding
objects with the Y − J < 0.5 that we applied to YJ-only selected
targets fainter than J = 18.5 since no T dwarfs have yet been found
with such colours. Although Y dwarfs are now known with such
blue Y − J colours (Leggett et al. 2013), we are not sensitive to such
objects due to their inherently faint nature.
One-third of our sample was selected using the new follow-up
method described in Section 3. Since this method uses the same
starting point of YJH and YJ-only colour selections, it is subject
to the same biases. In addition to those, however, it also includes
bias due to the z band and CH4 follow-up criteria. The follow-up
z -band observations delivered typically uncertainties of ±0.2 mag
for faint red T dwarfs, and considerably smaller uncertainties for
bluer M dwarfs. We ruled out from further investigation targets that
had z − J < 2.5 after z -band follow-up. This is roughly 3σ bluer
than the bluest measured z − J for late-T dwarfs. As such, we
estimate that less than 1 per cent of bona fide late-T dwarfs would
be scattered out of our selection during the z -band follow-up step.
Our CH4 photometry step may also exclude some objects
since we only obtained spectroscopy for those objects with
CH4 s − CH4 l < −0.5. The strong dependence of CH4 colour on
spectral type means that the fractional bias must be determined for
each spectral type bin. We have estimated the number of objects
that we would expect to scatter out of our CH4 s − CH4 l < −0.5 cut
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 15. A J-band-reduced proper motion versus spectral type diagram
for the UKIDSS late-T dwarfs with well measured proper motions. IsoVtan contours have been overplotted using by applying the Dupuy & Liu
(2012) MJ -spectral type relations. To apply these relations, we have adopted
Cushing et al. (2011) spectral types for the latest type objects.
by summing the probabilities of our confirmed dwarfs in each bin
being scattered beyond the cut based on their measured CH4 colours
and 1σ uncertainties. We find that we would expect 8 per cent of
T6–T6.5 dwarfs to be scattered out of our selection at this stage of
follow-up. Following the same method for T7–T7.5 dwarfs, we find
that less the 0.1 per cent will be excluded, and similarly negligible
fractions for the later type bins also.
We thus find that our bias correction factors for our full sample
due to our selection method are 1.03 for the T7–T7.5 and later bins,
and 1.05 for our T6–T6.5 bin. We must also correct for biases inherent to any magnitude-limited survey due to preferential inclusion of
unresolved binary systems and the Malmquist bias. In Burningham
et al. (2010b), we derived binary correction factors of 0.97 and 0.55
based on the broadest range of likely binary fractions reported in the
literature (5–45 per cent). In Pinfield et al. (2008), the Malmquist
correction factor was found to be between 0.84 and 0.86. We apply
these same correction factors here.
Calculating a space density from our corrected sample requires
the application of some MJ -spectral type conversion to determine
the volume probed for each subtype by our magnitude-limited sample. The choice of relation has a significant impact on the resulting space densities, and differences in approach to this issue can
contribute significantly to differences in measured space density
between different projects. In Burningham et al. (2010b), we used
the MJ -spectral type relations of Liu et al. (2006). However, to take
advantage of the improved sample of T dwarfs with measured parallaxes, we will instead use the polynomial relations presented by
Dupuy & Liu (2012) here. Since these relations are only valid for the
Cushing et al. (2011) system for the latest spectral types, we have
included our T9 dwarfs in the T8–T8.5 bin. In Table 10, we provide
a summary of our update to this calculation. By this method, we
find that the space density of T6–T8.5 dwarfs (on the Cushing et al.
2011 system) is 3.00 ± 1.3–5.52 ± 2.4 × 10−3 pc−3 , depending on
the underlying binary fraction.
Kirkpatrick et al. (2012) also used polynomial relations to estimate the magnitudes of objects for which no trigonometric parallax
is available and they applied a single 30 per cent unresolved binary
correction. To allow comparison between these samples, we have
carried out an additional space density calculation assuming the
same binary fraction. Table 11 summarizes this calculation, from
which we find a space density of 3.9 ± 1.7 × 10−3 pc−3 for T6–T8.5
dwarfs on the Cushing et al. (2011) system. This is very close to the
value of 3.43 ± 0.32 × 10−3 pc−3 found by Kirkpatrick et al. (2012).
The error we have assigned to Kirkpatrick value has been calculated
assuming only Poisson noise in their count of each spectral subtype.
Using polynomial relations is not necessarily, however, the most
appropriate way to estimate absolute magnitudes for targets of a
given subtype. This is because the subtypes are not a continuous
variable, but rather discrete and based on an inherently subjective
spectral typing scheme. As such, polynomial fits (and their associated residuals) can mask true scatter in the MJ values of certain
subtypes. We have thus performed an additional density estimate
using the mean magnitudes for each spectral subtype from Dupuy
& Liu (2012), and using the Cushing et al. (2011) spectral typing
scheme. For the purposes of this calculation, we have also used a
single binary correction factor of 30 per cent. Table 12 summarizes
this calculation that we include here for illustrative purposes.
This method leads us to an estimate of 6.2 ± 1.9 × 10−3 pc−3 for
the space density of T6–T8.5 dwarfs (on the system of Cushing et al.
2011). This is somewhat higher than the other methods determined
(for the same binary fraction), and it is driven by the fainter absolute
magnitude assigned to the coolest spectral type bin. Metchev et al.
(2008) also used a mean magnitude approach for estimating their
76 T dwarfs
489
Table 10. Summary of the updated space density calculation following the method of Burningham et al. (2010b) for our J < 18.8 sample of ≥T6 dwarfs. Nc
refers to corrected numbers based on the sample corrections described in the text, with maximum and minimum values arising from the different possible binary
corrections. The values of MJ used to calculate the distance limit and volume probed for each type were calculated using the polynomial relations in MJ versus
spectral type derived by Dupuy & Liu (2012). The uncertainties in MJ reflect the rms scatter about the Dupuy & Liu (2012) polynomials. The uncertainties
in the computed space densities reflect the volume uncertainty that arises from the uncertainty in MJ and Poisson noise in our sample. The minimum and
maximum space densities reflect the range encompassed by likely binary fractions (see the text and Burningham et al. 2010b). The latest spectral types are on
the Cushing et al. (2011) system.
Type
Teff range
N
Nc (min)
Nc (max)
MJ (MKO)
Range (pc)
Volume (pc3 )
ρ min (10−3 pc−3 )
ρ max (10−3 pc−3 )
T6–6.5
T7–7.5
T8–8.5
900–1050 K
800–900 K
500–800 K
40
21
17
19.4 ± 3.1
10.0 ± 2.2
6.7 ± 1.8
35.6 ± 5.7
18.5 ± 4.0
12.3 ± 3.3
14.90 ± 0.39
15.65 ± 0.39
16.75 ± 0.39
60 ± 11
43 ± 8
26 ± 5
50 400 ± 27 100
18 000 ± 9700
3900 ± 2100
0.39 ± 0.22
0.56 ± 0.32
2.05 ± 1.21
0.71 ± 0.40
1.02 ± 0.64
3.79 ± 2.24
Type
Teff range
N
Nc
MJ (MKO)
Range (pc)
Volume (pc3 )
ρ (10−3 pc−3 )
ρ (10−3 pc−3 )
Kirkpatrick et al. (2012)
T6–6.5
T7–7.5
T8–8.5
900–1050 K
800–900 K
500–800 K
40
21
17
25.3 ± 4.0
13.0 ± 2.8
10.5 ± 2.6
14.90 ± 0.39
15.65 ± 0.39
16.75 ± 0.39
60 ± 11
43 ± 8
26 ± 5
50400 ± 27100
18000 ± 9700
3900 ± 2100
0.50 ± 0.28
0.73 ± 0.42
2.67 ± 1.58
1.10 ± 0.18
0.93 ± 0.17
1.40 ± 0.21
Table 12. Summary of our space density calculation following the method of Burningham et al. (2010b) to correct biases
in our J < 18.8 sample of ≥T6 dwarfs. Nc refers to corrected numbers based on the sample corrections described in the
text, and a 30 per cent binary correction as applied by Kirkpatrick et al. (2012), and the latest spectral type bin is on the
system of Cushing et al. (2011). In this case, we have used the mean absolute magnitudes of Dupuy & Liu (2012) for our
MJ -spectral type conversion. We have applied these to each half subtype before combing them to arrive at space densities
for full subtype bins. The uncertainties in the computed space densities reflect the volume uncertainty that arises from
the uncertainty in MJ and Poisson noise in our sample.
Type
N
Nc
MJ (MKO)
Range (pc)
Volume (pc3 )
ρ (10−3 pc−3 )
T6
T6.5
T6–T6.5
T7
T7.5
T7–T7.5
T8
T8.5
T8–T8.5
24
16
40
14
7
21
12
5
17
15.2 ± 3.0
10.1 ± 2.5
15.22 ± 0.15
15.22 ± 0.31
52 ± 4
52 ± 7
32 400 ± 6700
32 400 ± 13 900
8.7 ± 2.3
4.3 ± 1.6
15.54 ± 0.25
16.05 ± 0.65
45 ± 5
35 ± 11
20 800 ± 7200
10 300 ± 9200
7.4 ± 2.1
3.1 ± 1.4
16.39 ± 0.35
17.81 ± 0.33
30 ± 5
16 ± 2
6400 ± 3100
900 ± 400
0.47 ± 0.12
0.31 ± 0.14
0.78 ± 0.19
0.42 ± 0.17
0.42 ± 0.39
0.84 ± 0.43
1.16 ± 0.63
3.43 ± 1.77
4.58 ± 1.88
survey volumes, applying a mean MJ (2MASS) = 15.75 ± 0.50 to
−3
pc−3
the T6–T8 bin, and deriving a space density of 4.3+2.9
−2.6 × 10
(for a binary fraction of ∼26 per cent). This is consistent with our
value of 2.8 ± 0.8 × 10−3 pc−3 across the same spectral type range.
Their slightly higher value can largely be attributed to fact that they
use a single mean for the entire T6–T8 bin, which is fainter than the
value we use for the most numerous T6 dwarfs. This will tend lead
to an overestimate of the space density across the whole bin. This
illustrates the significant impact that the assumed MJ -spectral-type
relation can have on derived space densities where trigonometric
parallaxes are not available for the entire sample. The uncertainties
introduced by the scatter about the assumed MJ , and the systematics
associated with the choice of MJ estimate relation, now significantly
outweigh the random error due to our sample size.
In Fig. 16, we compare our space densities from Table 11 and that
of Kirkpatrick et al. (2012), to the predictions from Monte Carlo
simulations of the Galactic field population of T dwarfs assuming
different forms of the IMF and brown dwarf formation history (see
Deacon & Hambly 2006; Burningham et al. 2010b, for full de-
scriptions). As in Burningham et al. (2010b) we have used the system mass function normalization of 0.0024 pc−3 for objects in the
0.09–0.10 M range, taken from Deacon, Nelemans & Hambly
(2008). For the purposes of this comparison we have transformed
the space densities which are supplied from the simulations as a
function of Teff , to densities as a function of spectral type using the
Teff ranges given in Table 11. These conversions are uncertain and
function as a very ‘broad brush’ to facilitate the comparison between
simulations and observation. We have also included predictions for
the Chabrier (2005) log-normal system mass function, since this
now appears to be the preferred function fitted to young clusters
across the low-mass stellar/substellar/planetary mass regime (e.g.
Bastian, Covey & Meyer 2010; Alves de Oliveira et al. 2012; Lodieu
et al. 2012b, and references therein).
It is clear that the estimated space densities from both UKIDSS
and WISE are significantly lower than predicted for an α = 0
or log-normal mass function, coinciding best with values of
−1.0 < α < −0.5. As has been noted before (e.g. Pinfield et al. 2008;
Burningham et al. 2010b), this represents a significant discrepancy
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Table 11. Summary of the updated space density calculation following the method of Burningham et al. (2010b) for our J < 18.8 sample of ≥T6 dwarfs. Nc
refers to corrected numbers based on the sample corrections described in the text, and 30 per cent binary correction as applied by Kirkpatrick et al. (2012). The
values of MJ and the uncertainties are as for Table 10. The latest spectral types are on the Cushing et al. (2011) system. The uncertainties for the Kirkpatrick
et al. (2012) densities estimated purely on the basis of Poisson noise on their number counts.
490
B. Burningham et al.
with the IMF measured in many young clusters where the mass
function has been fitted with α ≈ +0.6 power laws (e.g. Barrado y
Navascués et al. 2002; Lodieu et al. 2007c, 2009) or the log-normal
mass function with a characteristic mass in the region of 0.2 M that
has been measured by studies using the UKIDSS Galactic Clusters
Survey, amongst others (e.g. Lodieu et al. 2007a, 2009, 2011a,b,
2012b; Alves de Oliveira et al. 2012, 2013; Boudreault et al. 2012;
Lodieu, Deacon & Hambly 2012a). Kirkpatrick et al. (2012) note
that the ratio of stars to brown dwarfs in the field is consistent with
that seen in young clusters when a universal substellar IMF is assumed (Andersen et al. 2008), and thus conclude that the underlying
mass functions are the same. However, the discrepancy between the
predictions based on the functional form that has been fitted to the
IMF in young clusters, and observed field population should not be
ignored. Assuming a 30 per cent binary fraction, both our UKIDSS
T6–T8.5 space density and that of Kirkpatrick et al. (2012) are approximately half that predicted by the log-normal mass function,
and corresponding to a significance of 2σ and 13σ , respectively.
The kinematics of our sample (see Section 8) imply that this
discrepancy cannot be attributed to late-T dwarfs having a significantly different formation history to the rest of the Galactic disc,
and we are left with two plausible options. First, errors in the evolutionary models used to estimate masses in young clusters and/or
in the cooling times used to produce the Monte Carlo simulations
in field could give rise to such a discrepancy. These options can
be investigated by employing shifted cooling curves in the Monte
Carlo simulations to investigate their impact on the predicted space
densities and age profiles, and we defer further discussion of this
possibility to a future work.
The second plausible origin for such a discrepancy is if the bulk
of the field population formed in different environments than the
nearby young clusters that have been subject to detailed study. For
example, Bressert et al. (2010) have found that only 26 per cent
1 0 S U M M A RY
The expanded T dwarf sample presented here, along with access to
the new two epoch LAS catalogue of Smith et al. (in preparation) has
allowed the discovery of two new benchmark systems along with
the first estimate for the wide binary companion fraction amongst
late-T dwarfs. By examining the colours of the current census of
late-type benchmarks, we have identified the H − [4.5] and J −
W2 as being more sensitive to metallicity than models have so far
predicted, and caution that this must be considered when using these
colours to estimate the relative properties of cool brown dwarfs. The
expanded sample of late-T dwarfs, and the new model grids from
Saumon et al. (2012) and Morley et al. (2012) also now argue against
the previous conclusion of Leggett et al. (2010a) that the sample
is dominated by young low-mass objects. This is supported by the
kinematics of our sample.
Our updated space density for late-T dwarfs is consistent with
the density reported by Kirkpatrick et al. (2012) and confirms our
previous conclusion that there are far fewer late-T dwarfs than we
would expect given the favoured forms for the IMF in young clusters. Whilst it is possible that this discrepancy arises from problems
with either or both of the young and old evolutionary models for
brown dwarfs, we also speculate that it could arise as a result of
the dominant environment for low-mass star and brown dwarf formation being lower density regions than are currently probed, as
suggested for low-mass stars by Bressert et al. (2010).
AC K N OW L E D G E M E N T S
We thank our referee, J. Davy Kirkpatrick, for a helpful review
which substantially improved the quality of this manuscript. Based
on observations made under project A22TAC 96 on the Italian Telescopio Nazionale Galileo (TNG) operated on the island of La Palma
by the Fundacin Galileo Galilei of the INAF (Istituto Nazionale di
Astrofisica) at the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias. Based on observations obtained at the Gemini Observatory, which is operated by
the Association of Universities for Research in Astronomy, Inc.,
under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States),
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Figure 16. Computed space densities for different spectral types from
Monte Carlo simulations of the field population of T dwarfs for a uniform birthrate (i.e. β = 0.0, see Deacon & Hambly 2006; Burningham
et al. 2010b for full definition) and various underlying mass functions. The
power-law mass functions are of the form ψ(M) ∝ M−α pc−3 M−1
, and
the log-normal mass function is the Chabrier (2005) system mass function.
Our observed space density is shown as solid black lines, Kirkpatrick et al.
(2012) as dash–dotted black lines. Uncertainties are indicated with bars at
the mid-point of each spectral type bin, and reflect volume uncertainties and
Poisson counting uncertainties.
of low-mass stars form in high-density regions, which they define
as those with surface densities of young stellar objects (YSOs)
> 200 pc−2 , corresponding roughly to the most dense parts of the
Taurus star-forming complex. Luhman (2004) found a deficit of
brown dwarfs in Taurus relative to the high densities of the Trapezium cluster, and more recent studies have further hinted at similar enrichment of the substellar component at higher densities
(Andersen et al. 2011). A combination of low-density-dominated
low-mass star and brown dwarf formation with a substellar IMF
that declines steeply in the lowest density regions, but has the lognormal form currently preferred in higher density regions, would
naturally reconcile the two conflicting views of form for the substellar IMF that we are currently presenting. Unfortunately, obtaining
a statistically meaningful census of the young substellar population
in the lowest density regions where Bressert et al. (2010) argue that
most low-mass stars are formed (e.g. <50 YSOs pc−2 ) is extremely
challenging due to the large areas that must be surveyed and the inherently faint nature of the targets. Given the significant investment
of observing time that would be required to test this hypothesis, it
would make sense to first investigate the potential impact of systematic effects in the evolutionary models for both young and old
objects.
76 T dwarfs
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the Science and Technology Facilities Council (United Kingdom),
the National Research Council (Canada), CONICYT (Chile), the
Australian Research Council (Australia), Ministério da Ciência e
Tecnologia (Brazil) and Ministerio de Ciencia, Tecnologı́a e Innovación Productiva (Argentina). We would like to acknowledge
the support of the Marie Curie 7th European Community Framework Programme grant no. 247593 Interpretation and Parametrization of Extremely Red COOL dwarfs (IPERCOOL) International
Research Staff Exchange Scheme. ADJ is supported by a Fondecyt postdoctorado fellowship, under project number 3100098, and
is also partially supported by the proyecto Basal PB06 (CATA)
and the Joint Committee ESO-Government Chile. AHA thanks
CNPq grant PQ306775/2009-3 and SHAO/CAS Visiting Professorship grant. CGT is supported by ARC grant DP0774000. SKL’s
research is supported by the Gemini Observatory. JG is supported
by RoPACS, a Marie Curie Initial Training Network funded by the
European Commission’s Seventh Framework Programme. NL acknowledges funding from the Spanish Ministry of Science and Innovation through the Ramón y Cajal fellowship number 08-303-01-02
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summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to
conduct observations from this mountain.
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A P P E N D I X A : S Q L DATA BA S E QU E R I E S U S E D
F O R IN I T I A L C A N D I DAT E S E L E C T I O N
The following code listings give the SQL code we used to query
the WSA for our candidates. These include a cross-match against
SDSS DR7. To access SDSS DR8 sky not covered by SDSS DR7
we independently cross-matched candidate lists that had passed our
NIR colour cuts against SDSS DR8. Our final J − K < 0.1 colour
cut was applied to the resulting candidate lists.
Query 1. The SQL query used to select candidates via our YJH
channel in the case where they have been detected in SDSS.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Leggett S. K., Saumon D., Burningham B., Cushing M. C., Marley M. S.,
Pinfield D. J., 2010b, ApJ, 720, 252
Leggett S. K., Morley C. V., Marley M. S., Saumon D., Fortney J. J., Visscher
C., 2013, ApJ, 763, 130
Lepine S., Hilton E. J., Mann A. W., Wilde M., Rojas-Ayala B., Cruz K. L.,
Gaidos E., 2013, AJ, 145, 102
Lépine S., Shara M. M., 2005, AJ, 129, 1483
Liu M. C., Leggett S. K., Golimowski D. A., Chiu K., Fan X., Geballe
T. R., Schneider D. P., Brinkmann J., 2006, ApJ, 647, 1393
Liu M. C., Leggett S. K., Chiu K., 2007, ApJ, 660, 1507
Liu M. C., Dupuy T. J., Bowler B. P., Leggett S. K., Best W. M. J., 2012,
ApJ, 758, 57
Lodders K., Fegley B., 2002, Icarus, 155, 393
Lodieu N., Hambly N. C., Jameson R. F., Hodgkin S. T., Carraro G., Kendall
T. R., 2007a, MNRAS, 374, 372
Lodieu N. et al., 2007b, MNRAS, 379, 1423
Lodieu N., Dobbie P. D., Deacon N. R., Hodgkin S. T., Hambly N. C.,
Jameson R. F., 2007c, MNRAS, 380, 712
Lodieu N., Zapatero Osorio M. R., Rebolo R., Martı́n E. L., Hambly N. C.,
2009, A&A, 505, 1115
Lodieu N., Hambly N. C., Dobbie P. D., Cross N. J. G., Christensen L.,
Martin E. L., Valdivielso L., 2011a, MNRAS, 418, 2604
Lodieu N., de Wit W.-J., Carraro G., Moraux E., Bouvier J., Hambly N. C.,
2011b, A&A, 532, A103
Lodieu N., Deacon N. R., Hambly N. C., 2012a, MNRAS, 422, 1495
Lodieu N., Deacon N. R., Hambly N. C., Boudreault S., 2012b, MNRAS,
426, 3403
Lodieu N. et al., 2012c, A&A, 548, A53
Looper D. L., Kirkpatrick J. D., Burgasser A. J., 2007, AJ, 134, 1162
Luhman K. L., 2004, ApJ, 617, 1216
Luhman K. L. et al., 2012, ApJ, 760, 152
Mace G. N. et al., 2013, ApJS, 205, 6
Manchado A. et al., 1998, in Fowler A. M., ed., Proc. SPIE Vol. 3354,
Infrared Astronomical Instrumentation. SPIE, Bellingham, p. 448
Mann A. W., Brewer J. M., Gaidos E., Lepine S., Hilton E. J., 2013, AJ,
145, 52
Marocco F. et al., 2010, A&A, 524, A38
Masana E., Jordi C., Ribas I., 2006, A&A, 450, 735
Metchev S. A., Kirkpatrick J. D., Berriman G. B., Looper D., 2008, ApJ,
676, 1281
Molinari E., Conconi P., Pucillo M., 1997, Mem. Soc. Astron. Ital., 68, 231
Morley C. V., Fortney J. J., Marley M. S., Visscher C., Saumon D., Leggett
S. K., 2012, ApJ, 756, 172
Murray D. N. et al., 2011, MNRAS, 414, 575
Neves V. et al., 2012, A&A, 538, A25
Oke J. B., 1990, AJ, 99, 1621
Perryman M. A. C. et al., 1997, A&A, 323, L49
Pinfield D. J., Jones H. R. A., Lucas P. W., Kendall T. R., Folkes S. L.,
Day-Jones A. C., Chappelle R. J., Steele I. A., 2006, MNRAS, 368,
1281
Pinfield D. J. et al., 2008, MNRAS, 390, 304
Pinfield D. J. et al., 2012, MNRAS, 422, 1922
Roche P. F. et al., 2003, in Iye M., Moorwood A. F. M., eds, Proc. SPIE
Vol. 4841, Instrument Design and Performance for Optical/Infrared
Ground-based Telescopes. SPIE, Bellingham, p. 901
Rojas-Ayala B., Covey K. R., Muirhead P. S., Lloyd J. P., 2012, ApJ, 748,
93
Saumon D., Marley M. S., 2008, ApJ, 689, 1327
Saumon D., Marley M. S., Cushing M. C., Leggett S. K., Roellig T. L.,
Lodders K., Freedman R. S., 2006, ApJ, 647, 552
Saumon D., Marley M. S., Abel M., Frommhold L., Freedman R. S., 2012,
ApJ, 750, 74
Schlaufman K. C., Laughlin G., 2010, A&A, 519, A105
Scholz R., 2010, A&A, 515, 92
Scholz R.-D., Bihain G., Schnurr O., Storm J., 2012, A&A, 541, A163
Simcoe R. A. et al., 2008, in McLean I. S., Casali M. M., eds, Proc. SPIE
Vol. 7014, Ground-based and Airborne Instrumentation for Astronomy
II. SPIE, Bellingham, p. 70140U
76 T dwarfs
493
Query 2. The SQL query used to select candidates via our YJH
channel in the case where they have not been detected in SDSS.
Query 4. The SQL query used to select candidates via our YJ-only
channel in the case where they have not been detected in SDSS.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
Query 3. The SQL query used to select candidates via our YJ-only
channel in the case where they have been detected in SDSS.
SELECT s . ra , s . dec , s . ymj 1Pnt , s . ym j 1P n t E r r , s . j 1mhPnt
, s . j 1m h P n t E r r ,
s . yAperMag3 , s . yAperMag3Err , s . j 1AperMag3 , s .
j 1A p er Mag3E rr , my . mjdObs − mj . mjdObs , s .
mergedClass
/∗ O n ly f r a m e s w i t h f u l l c o v e r a g e ∗/
FROM
lasYJHKSource As s , lasYJHKmergeLog AS l ,
M u l t i f r a m e AS my, M u l t i f r a m e AS mj
WHERE
/∗
C o l o u r c u t s f o r mid−T & l a t e r o r b r i g h t e n o u g h
t h a t i t s n o t M d w a r f : ∗/
( s . ymj 1Pnt > 0 . 5
OR s . j 1A p erMag3 < 1 8 . 5 )
AND
/∗
U n d u p l i c a t e d o r p r i m a r y d u p l i c a t e s o n l y : ∗/
( s . p r i O r S e c = 0 OR s . p r i O r S e c = s . f r a m e S e t I D )
AND
/∗
G e n e r a l l y g o o d q u a l i t y : ∗/
s . yppErrBits
< 256 AND
s . j 1 p p E r r B i t s < 256 AND
s . f r a m e S e t I D = l . f r a m e S e t I D AND
/∗
P i c k o u t t h e YJHK f r a m e s t o g e t t h e m jds ∗/
l . ymfID = my . m u l t i f r a m e I D AND
l . j 1 m f I D = mj . m u l t i f r a m e I D AND
/∗
S o u r c e n o t d e t e c t e d a b o v e 2 s i g m a i n SDSS−DR7 i ’
o r z ’ w i t h i n 2 a r c s e c : ∗/
s o u r c e I D NOT IN (
SELECT m aster O b jID
FROM
lasSourceXDR7PhotoObj AS x ,
BestDR7 . . PhotoObj
AS p
WHERE p . ob jID = x . s l a v e O b j I D
AND
( p s f M a g E r r i < 0 . 5 OR p s f M a g E r r z < 0 . 5 )
AND
x . distanceMins < 1.0/60.0
) AND
/∗
Use o n l y f r a m e s e t s o v e r l a p p i n g w i t h SDSS−DR7 :
∗/
s . f r a m e S e t I D IN (
SELECT DISTINCT( s . f r a m e S e t I D )
FROM
lasSource
AS n ,
lasSourceXDR7PhotoObj AS x
WHERE n . s o u r c e I D = x . m aster O b jID
) AND
/∗
S t a r − l i k e m o r p h o l o g i c a l c l a s s i f i c a t i o n : ∗/
s . m e r g e d C l a s s BETWEEN −2.0 and −1.0 AND
s . m e r g e d C l a s s S t a t BETWEEN −3.0 AND +3.0 AND
/∗
R e a s o n a b l y c i r c u l a r i m a g e s : ∗/
s . y E l l < 0 . 4 5 AND s . j 1 E l l < 0 . 4 5 AND
/∗
IR p a i r s w i t h i n 0 . 7 5 a r c s e c : ∗/
( ( ( s . yXi BETWEEN −0.75 AND + 0 . 7 5 ) AND
( s . yEta
BETWEEN −0.75 AND + 0 . 7 5 ) AND
( s . j 1 X i BETWEEN −0.75 AND + 0 . 7 5 ) AND
( s . j 1 E t a BETWEEN −0.75 AND + 0 . 7 5 ) ) OR
/∗
Or MJD o f OBS s e p a r a t e d b y more t h a n 1 d a y s ∗/
( ( my . mjdObs − mj . mjdObs ) > 1 or ( mj . mjdObs − my .
mjdObs) > 1 ) ) AND
/∗
YJ m e a s u r e d t o 3 : ∗/
s . yAperMag3Err
< 0 . 3 0 AND s . yAperMag3Err
>
0 AND
s . j 1A p er Mag3E rr < 0 . 3 0 AND s . j 1A p er Mag3E rr >
0 AND
/∗ H and K d r o p o u t : ∗/
s . hAPerMag3Err < 0 AND s . kAperMag3Err < 0 AND
s . hAperMag3 < 0 AND s . kAperMag3 < 0 AND
/∗
J b r i g h t e r t h a n 1 9 . 5 ∗/
s . j 1A p er Mag3 < 1 9 . 3 0
ORDER BY r a
494
B. Burningham et al.
A P P E N D I X B : S U M M A RY O F P H OT O M E T R I C
O B S E RVAT I O N S
Table B1. Summary of broad-band photometric observations.
Object
ULAS J0128+0633
ULAS J0130+0804
ULAS J0226+0702
ULAS J0329+0430
ULAS J0746+2355
ULAS J0747+2455
ULAS J0758+2225
ULAS J0759+1855
ULAS J0800+1908
ULAS J0809+2126
ULAS J0814+2452
ULAS J0815+2711
ULAS J0821+2509
ULAS J0847+0350
ULAS J0927+3413
ULAS J0929+0409
ULAS J0950+0117
ULAS J0954+0623
ULAS J1021+0544
Instrument
z
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO J
MKO J
MKO H
MKO Y
MKO J
MKO H
MKO K
z
H
CH4 l
MKO Y
MKO J
MKO H
MKO K
MKO J
MKO H
z
z
MKO J
MKO H
z
H
CH4 l
MKO J
MKO H
MKO J
MKO H
z
MKO Y
MKO J
MKO H
MKO K
z
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
ACAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
LIRIS
LIRIS
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
DOLORES
WFCAM
WFCAM
DOLORES
LIRIS
LIRIS
WFCAM
WFCAM
WFCAM
WFCAM
ACAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
UT
date
2011-01-08
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2009-12-16
2009-12-16
2009-12-16
2010-11-24
2010-11-24
2010-11-24
2010-11-24
2010-12-27
2011-01-09
2011-01-09
2010-04-19
2010-04-19
2010-04-19
2010-04-19
2009-12-16
2009-12-16
2010-12-18
2010-12-28
2010-01-08
2010-01-08
2010-12-28
2011-01-09
2011-01-09
2010-01-10
2010-01-10
2010-01-11
2010-01-11
2011-01-12
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2011-05-11
2010-11-22
2010-11-22
2010-11-22
2010-11-22
2010-01-08
2009-12-16
2009-12-16
2010-01-08
2010-11-23
2010-11-23
2010-11-23
2010-11-23
2010-11-25
2010-11-25
2010-11-25
2010-11-25
Programme ID
Tint (s)
W/10B/P16
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/09B/7
U/09B/7
U/09B/7
U/10B/8
U/10B/8
U/10B/8
U/10B/8
A22TAC_96
W/10B/P16
W/10B/P16
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/09B/7
U/09B/7
A22TAC_96
A22TAC_96
U/09B/7
U/09B/7
A22TAC_96
W/2010B/P16
W/2010B/P16
U/09B/7
U/09B/7
U/09B/7
U/09B/7
W/10B/P16
U/10B/8
U/10B/8
U/10B/8
U/10B/8
A23TAC_28
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/09B/7
U/09B/7
U/09B/7
U/09B/7
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
1800
280
120
1000
1000
280
120
1000
1000
120
120
1000
280
120
1000
1000
1200
3000
3000
280
120
1000
1000
120
1000
1200
1200
120
1000
1200
740
740
120
1000
120
1000
1200
280
120
1000
1000
1200
280
120
1000
1000
280
120
1000
1000
280
120
1000
1000
280
120
1000
1000
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
ULAS J0245+0653
ULAS J0255+0616
Filter
76 T dwarfs
495
Table B1 – continued
Object
ULAS J1023+0447
ULAS J1029+0935
ULAS J1043+1048
ULAS J1137+1126
ULAS J1204−0150
ULAS J1206+1018
ULAS J1212+1010
ULAS J1258+0307
ULAS J1302+1434
ULAS J1335+1506
ULAS J1338−0142
ULAS J1339−0056
ULAS J1421+0136
ULAS J1425+0451
ULAS J1516+0110
ULAS J1534+0556
ULAS J1549+2621
ULAS J1601+2646
ULAS J1614+2442
ULAS J1619+2358
ULAS J1619+3007
ULAS J1626+2524
ULAS J2116−0101
Instrument
z
MKO Y
MKO J
H
CH4 l
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
z
MKO J
MKO H
MKO J
MKO H
MKO J
MKO H
MKO J
MKO H
z
MKO J
MKO H
MKO Y
MKO J
MKO H
MKO Y
MKO J
MKO H
MKO K
MKO Y
MKO J
MKO H
MKO K
MKO J
MKO H
z
MKO J
MKO H
z
MKO J
MKO H
z
MKO Y
MKO J
MKO H
MKO K
z
MKO J
MKO H
MKO J
MKO H
MKO J
MKO H
z
DOLORES
WFCAM
WFCAM
LIRIS
LIRIS
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
UFTI
UFTI
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
UFTI
UFTI
UFTI
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
UT
date
2010-12-28
2010-11-26
2010-11-26
2011-01-09
2011-01-09
2010-11-26
2010-11-26
2010-11-23
2010-11-23
2010-11-23
2010-11-23
2010-12-06
2010-12-06
2010-12-06
2010-12-06
2010-12-06
2010-12-06
2010-12-06
2010-12-06
2011-05-08
2009-01-25
2009-01-25
2009-07-14
2009-07-14
2009-07-13
2009-07-13
2010-04-20
2010-04-20
2011-07-07
2009-07-12
2009-07-12
2009-01-11
2009-01-11
2009-01-11
2010-05-13
2010-05-13
2010-05-13
2010-05-13
2010-05-12
2010-05-12
2010-05-12
2010-05-12
2009-07-12
2009-07-12
2011-05-12
2010-04-20
2010-04-20
2011-05-12
2010-05-09
2010-05-09
2011-05-14
2010-05-15
2010-05-15
2010-05-15
2010-05-15
2011-05-09
2010-05-11
2010-05-11
2010-05-12
2010-05-12
2010-06-13
2010-06-13
2011-05-11
Programme ID
Tint (s)
A22TAC_96
U/10B/8
U/10B/8
W/2010B/P16
W/2010B/P16
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
U/10B/8
A23TAC_28
U/08B/15
U/08B/15
U/09A/1
U/09A/1
U/09A/1
U/09A/1
U/10A/6
U/10A/6
A23TAC_28
U/09A/1
U/09A/1
U/08B/15
U/08B/15
U/08B/15
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/09A/1
U/09A/1
A23TAC_28
U/10A/6
U/10A/6
A23TAC_28
U/10A/6
U/10A/6
A23TAC_28
U/10A/6
U/10A/6
U/10A/6
U/10A/6
A23TAC_28
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
U/10A/6
A23TAC_28
900
280
120
900
900
1000
1000
280
120
1000
1000
280
120
1000
1000
280
120
1000
1000
900
300
1800
120
1000
120
1000
120
1000
1200
120
1000
540
300
1800
280
120
1000
1000
280
120
1000
1000
120
1000
1200
120
1000
1800
120
1000
1650
280
120
1000
1000
1200
120
1000
120
1000
120
1000
900
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
ULAS J1051−0154
Filter
496
B. Burningham et al.
Table B1 – continued
Object
ULAS J2237+0642
ULAS J2300+0703
ULAS J2357+0132
Instrument
MKO Y
MKO J
MKO H
MKO K
z
MKO Y
MKO J
MKO H
MKO K
z
z
z
MKO Y
MKO J
MKO H
MKO K
z
MKO Y
MKO J
MKO H
MKO K
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
DOLORES
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
DOLORES
WFCAM
WFCAM
WFCAM
WFCAM
UT
date
2010-07-11
2010-11-22
2010-11-22
2010-11-03
2011-07-08
2009-07-14
2009-07-14
2009-07-14
2009-07-14
2011-11-16
2012-01-13
2011-11-16
2010-11-25
2010-11-25
2010-11-25
2010-11-25
2011-07-08
2010-11-22
2010-11-22
2010-11-22
2010-11-22
Programme ID
Tint (s)
U/10B/8
U/10B/8
U/10B/8
U/10B/8
A23TAC_28
U/09A/1
U/09A/1
U/09A/1
U/09A/1
A24TAC_49
A24TAC_49
A24TAC_49
U/10B/8
U/10B/8
U/10B/8
U/10B/8
A23TAC_28
U/10B/8
U/10B/8
U/10B/8
U/10B/8
280
260
2000
2000
900
120
120
400
400
900
900
900
280
120
1000
1000
900
280
120
1000
1000
Table B2. Summary of methane photometric observations using the TNG. Tint gives the
integration time for each co-add, Ncoadd is the number of co-added images at each dither
point, Ndither is the number of dither points in the mosaic.
Name
ULAS J0007+0112
ULAS J0127+1539
ULAS J0128+0633
ULAS J0130+0804
ULAS J0133+0231
ULAS J0139+1503
ULAS J0200+0908
ULAS J0745+2332
ULAS J0759+1855
ULAS J0811+2529
ULAS J0847+0350
ULAS J0926+0402
ULAS J0927+3413
ULAS J0929+0409
ULAS J0954+0623
ULAS J1021+0544
ULAS J1029+0935
ULAS J1042+1212
ULAS J1043+1048
ULAS J1051−0154
ULAS J1053+0157
ULAS J1137+1126
ULAS J1152+0359
ULAS J1152+1134
ULAS J1223-0131
ULAS J1228+0407
ULAS J1254+1222
ULAS J1259+2933
ULAS J1302+1434
ULAS J1335+1506
ULAS J1417+1330
ULAS J1425+0451
ULAS J1449+1147
ULAS J1516+0110
TNG programme
A24TAC_49
A24TAC_49
A22TAC_96
A22TAC_96
A24TAC_49
A24TAC_49
A24TAC_49
A24TAC_49
A24TAC_49
A22TAC_96
A23TAC_28
A25TAC_32
A23TAC_28
A23TAC_28
A23TAC_28
A24TAC_49
A23TAC_28
A24TAC_49
A23TAC_28
A23TAC_28
A24TAC_49
A23TAC_28
A22TAC_96
A24TAC_49
A25TAC_32
A23TAC_28
A24TAC_49
A24TAC_49
A24TAC_49
A23TAC_28
A23TAC_28
A23TAC_28
A23TAC_28
A23TAC_28
UT
date
2011-10-27
2011-11-19
2010-11-06
2010-12-25
2011-11-18
2011-10-27
2011-10-28
2011-10-28
2011-10-27
2010-12-27
2011-05-07
2012-04-29
2011-05-12
2011-05-07
2011-05-09
2012-01-16
2011-05-09
2012-01-16
2011-05-09
2011-05-10
2012-01-16
2011-05-10
2010-12-26
2012-01-17
2012-04-29
2011-05-11
2012-01-14
2012-02-01
2012-01-16
2011-05-10
2011-05-07
2011-05-13
2011-05-07
2011-07-09
Tint (s)
Ncoadds
Ndither
30.0
20.0
20.0
30.0
20.0
45.0
30.0
60.0
60.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
20.0
30.0
30.0
20.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
30.0
1
2
3
1
2
1
1
1
1
1
1
4
2
1
1
4
1
4
2
1
4
2
1
4
4
1
4
4
4
1
1
2
1
1
30
30
30
30
30
30
30
30
30
30
30
10
30
30
30
10
30
10
30
30
10
30
30
10
10
30
10
10
10
30
30
30
30
30
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
ULAS J2318+0433
ULAS J2326+0509
ULAS J2331+0426
ULAS J2342+0856
Filter
76 T dwarfs
497
Table B2 – continued
Name
TNG programme
A23TAC_28
A23TAC_28
A23TAC_28
A23TAC_28
A23TAC_28
A23TAC_28
A24TAC_49
A23TAC_28
A24TAC_49
A24TAC_49
A24TAC_49
A24TAC_49
A24TAC_49
Table C1. Dates, instruments and programme numbers for spectra obtained for this work.
Instrument
ULAS J0007+0112
ULAS J0127+1539
ULAS J0128+0633
ULAS J0130+0804
ULAS J0133+0231
ULAS J0139+1503
ULAS J0200+0908
ULAS J0226+0702
ULAS J0245+0653
ULAS J0255+0616
ULAS J0329+0430
ULAS J0745+2332
ULAS J0746+2355
ULAS J0747+2455
ULAS J0758+2225
ULAS J0759+1855
ULAS J0809+2126
ULAS J0811+2529
ULAS J0814+2452
ULAS J0815+2711
ULAS J0819+2103
ULAS J0821+2509
ULAS J0926+0402
ULAS J0927+2524
ULAS J0929+0409
ULAS J0932+3102
ULAS J0950+0117
ULAS J0950+0117
ULAS J0950+0117
ULAS J0954+2452
ULAS J1021+0544
ULAS J1023+0447
ULAS J1029+0935
ULAS J1042+1212
ULAS J1043+1048
ULAS J1051+0154
ULAS J1053+0157
ULAS J1111+0518
ULAS J1152+0359
ULAS J1152+1134
GNIRS
GNIRS
GNIRS
GNIRS
GNIRS
GNIRS
GNIRS
GNIRS
IRCS
XSHOOTER
NIRI
GNIRS
GNIRS
IRCS
IRCS
GNIRS
IRCS
GNIRS
GNIRS
IRCS
NIRI
NIRI
FIRE
GNIRS
IRCS
NIRI
IRCS
NIRI(H)
NIRI(K)
FIRE
IRCS
GNIRS
IRCS
GNIRS
GNIRS
IRCS
FIRE
FIRE
IRCS
GNIRS
UT
date
2011-11-14
2011-12-31
2010-12-15
2011-10-15
2011-12-17
2011-11-24
2011-11-26
2011-10-15
2009-12-30
2010-12-01
2009-11-03
2011-11-22
2011-04-23
2009-12-30
2009-12-31
2011-11-25
2009-12-30
2010-12-30
2011-04-19
2009-12-30
2009-11-01
2009-12-31
2012-05-09
2011-06-16
2011-01-24
2009-12-31
2009-05-07
2009-12-08
2009-12-30
2012-05-09
2011-01-23
2011-06-12
2011-01-23
2012-03-05
2011-06-17
2011-01-24
2012-05-09
2012-05-09
2011-01-24
2012-06-06
2011-05-13
2011-07-10
2011-05-10
2011-05-07
2011-07-11
2011-05-08
2011-10-26
2011-07-09
2011-10-27
2012-01-16
2011-10-27
2012-01-16
2011-10-28
Tint (s)
Ncoadds
Ndither
30.0
30.0
30.0
30.0
30.0
30.0
45.0
30.0
60.0
30.0
45.0
30.0
30.0
2
1
1
1
2
2
1
1
1
4
1
4
1
30
30
30
30
30
30
30
30
30
10
30
10
30
Table C1 – continued
A P P E N D I X C : S U M M A RY O F
S P E C T RO S C O P I C O B S E RVAT I O N S
Target
date
Programme ID
GN-2011B-Q-5
GN-2011B-Q-43
GN-2010B-Q-41
GN-2011B-Q-5
GN-2011B-Q-43
GN-2011B-Q-43
GN-2011B-Q-43
GN-2011B-Q-5
o09164
086.C-0450(A)
GN-2009B-Q-62
GN-2011B-Q-43
GN-2011A-Q-73
o09164
o09164
GN-2011B-Q-5
o09164
GN-2010B-Q-41
GN-2011A-Q-73
o09164
GN-2009B-Q-62
GN-2009B-Q-62
GN-2011A-Q-73
o10148
GN-2009B-Q-62
o09118
GN-2009B-Q-62
GN-2009B-Q-62
o10148
GN-2011A-Q-73
o10148
GN-2012A-Q-84
GN-2011A-Q-73
o10148
o10148
GN-2012A-Q-84
Target
Instrument
UT Date
Programme ID
ULAS J1155+0445
ULAS J1204+0150
ULAS J1206+1018
ULAS J1212+1010
ULAS J1223−0131
ULAS J1228+0407
ULAS J1258+0307
ULAS J1259+2933
ULAS J1302+1434
ULAS J1335+1506
ULAS J1338−0142
ULAS J1339−0056
ULAS J1339+0104
ULAS J1417+1330
ULAS J1421+0136
ULAS J1425+0451
ULAS J1449+1147
ULAS J1516+0110
ULAS J1517+0529
ULAS J1534+0556
ULAS J1536+0155
ULAS J1549+2621
ULAS J1601+2646
ULAS J1614+2442
ULAS J1617+2350
ULAS J1619+2358
ULAS J1619+3007
ULAS J1626+2524
ULAS J1639+3232
ULAS J2116−0101
ULAS J2237+0642
ULAS J2300+0703
ULAS J2315+0344
ULAS J2318+0433
ULAS J2326+0201
ULAS J2331+0426
ULAS J2342+0856
ULAS J2352+1244
ULAS J2357+0132
NIRI
NIRI
NIRI
NIRI
FIRE
FIRE
IRCS
GNIRS
FIRE
IRCS
NIRI
IRCS
IRCS
GNIRS
NIRI
GNIRS
GNIRS
IRCS
GNIRS
GNIRS
IRCS
IRCS
GNIRS
GNIRS
GNIRS
NIRI
GNIRS
IRCS
NIRI
GNIRS
GNIRS
NIRI
GNIRS
GNIRS
GNIRS
GNIRS
NIRI
GNIRS
GNIRS
2009-12-31
2009-04-16
2010-02-06
2010-01-28
2012-05-09
2012-05-09
2010-04-05
2012-04-20
2012-05-09
2009-05-06
2010-05-01
2010-04-05
2010-04-05
2011-05-16
2010-01-28
2011-07-09
2011-05-15
2010-04-06
2011-08-13
2011-07-09
2010-04-05
2010-04-06
2011-03-17
2011-07-10
2011-05-15
2010-04-27
2012-05-31
2010-04-05
2010-04-30
2011-11-14
2010-12-07
2009-06-15
2011-12-24
2012-06-02
2011-11-28
2012-06-06
2009-10-31
2011-11-16
2010-12-08
GN-2009B-Q-62
GN-2009A-Q-16
GN-2010A-Q-44
GN-2009B-Q-62
o10121
GN-2012A-Q-84
o09118
GN-2010A-Q-44
o10121
o10121
GN-2011A-Q-73
GN-2009B-Q-62
GN-2011A-Q-73
GN-2011A-Q-73
o10121
GN-2011B-Q-5
GN-2011A-Q-73
o10121
o10121
GN-2011A-Q-73
GN-2011A-Q-73
GN-2011A-Q-73
GN-2010A-Q-44
GN-2012A-Q-84
o10121
GN-2010A-Q-44
GN-2011B-Q-5
GN-2010B-Q-41
GN-2009A-Q-16
GN-2011B-Q-43
GN-2012A-Q-84
GN-2011B-Q-43
GN-2012A-Q-84
GN-2009B-Q-62
GN-2011B-Q-43
GN-2010B-Q-41
This paper has been typeset from a TEX/LATEX file prepared by the author.
Downloaded from http://mnras.oxfordjournals.org/ at Isaac Newton Group of Telescopes on April 7, 2014
ULAS J1534+0556
ULAS J1549+2621
ULAS J1614+2442
ULAS J1617+2350
ULAS J1619+2358
ULAS J1619+3007
ULAS J2116−0101
ULAS J2300+0703
ULAS J2315+0344
ULAS J2318+0433
ULAS J2326+0201
ULAS J2326+0509
ULAS J2352+1244
UT
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