+ white dwarf binary system Discovery of a T dwarf

+ white dwarf binary system Discovery of a T dwarf
Mon. Not. R. Astron. Soc. 410, 705–716 (2011)
Discovery of a T dwarf + white dwarf binary system
A. C. Day-Jones,1,2 † D. J. Pinfield,2 M. T. Ruiz,1 H. Beaumont,2 B. Burningham,2
J. Gallardo,1 A. Gianninas,3 P. Bergeron,3 R. Napiwotzki,2 J. S. Jenkins,1
Z. H. Zhang,2 D. N. Murray,2 S. Catalán2 and J. Gomes2
1 Departamento
de Astronomia, Universidad de Chile, Camino del Observatorio 1515, Santiago, Chile
for Astrophysics Research, University of Hertfordshire, College Lane, Hatfield, Hertfordshire AL10 9AB
3 Département de Physique, Université de Montréal, CP 6128, Succursale Centre-Ville, Montreal, Canada
2 Centre
Accepted 2010 August 3. Received 2010 August 2; in original form 2010 March 16
We present the discovery of the first T dwarf + white dwarf binary system
LSPM 1459+0857 AB, confirmed through common proper motion and spectroscopy. The
white dwarf is a high proper motion object from the LSPM catalogue that we confirm spectroscopically to be a relatively cool (T eff = 5535 ± 45 K) and magnetic (B ∼ 2 MG) hydrogen-rich
white dwarf, with an age of at least 4.8 Gyr. The T dwarf is a recent discovery from the UKIRT
Infrared Deep Sky Survey (ULAS 1459+0857) and has a spectral type of T4.5 ± 0.5 and
a distance in the range 43–69 pc. With an age constraint (inferred from the white dwarf)
of >4.8 Gyr, we estimate T eff = 1200–1500 K and log g = 5.4–5.5 for ULAS 1459+0857,
making it a benchmark T dwarf with well-constrained surface gravity. We also compare the T
dwarf spectra with the latest LYON group atmospheric model predictions, which, despite some
shortcomings, are in general agreement with the observed properties of ULAS 1459+0857.
The separation of the binary components (16 500–26 500 au or 365 arcsec on the sky) is consistent with an evolved version of the more common brown dwarf + main-sequence (MS)
binary systems now known, and although the system has a wide separation, it is shown to be
statistically robust as a non-spurious association. The observed colours of the T dwarf show
that it is relatively bright in the z band compared to other T dwarfs of similar type, and further
investigation is warranted to explore the possibility that this could be a more generic indicator of older T dwarfs. Future observations of this binary system will provide even stronger
constraints on the T dwarf properties, and additional systems will combine to give a more
comprehensively robust test of the model atmospheres in this temperature regime.
Key words: binaries: general – brown dwarfs – white dwarfs.
1 I N T RO D U C T I O N
Large-scale near-infrared (NIR) and optical surveys, such as
the Two-Micron All-Sky Survey (2MASS), the Sloan Digital
Sky Survey (SDSS) and the UKIRT Infrared Deep Sky Survey
(UKIDSS), are aiding the identification of a rapidly increasing
number of ‘field’ brown dwarfs (BDs) (e.g. Lodieu et al. 2007;
Pinfield et al. 2008; Burningham et al. 2009), as well as probing
down into new cooler temperature regimes (Warren et al. 2007;
Burningham et al. 2008; Delorme et al. 2008; Burningham et al.
E-mail: [email protected]
†Based on observations made with ESO telescopes at the La Silla Paranal
Observatory under programme 282.C-5069(A).
C 2010 RAS
2010 The Authors. Journal compilation 2009; Leggett et al. 2009). In general, the estimation of properties
of these BDs (e.g. age, mass, metallicity) currently relies on model
fitting. However, the models are very sensitive to a variety of poorly
understood processes in BD atmospheres, such as the formation of
dust condensates (Allard et al. 2001) and non-equilibrium chemistry
(Saumon et al. 2007) and the spectroscopic fitting of atmospheric
properties (T eff , log g, [M/H]) is a major challenge. Crucially, the
nature of BD evolution means that the mass–luminosity relation
depends strongly on the age, and in the absence of well-constrained
atmospheric properties, there is no way to accurately determine the
mass and age.
Identifying objects, where one can pin down these properties
independently, can help aid the calibration of models. BDs that
are members of binaries, where the primary member has age constraints, are good sources of benchmark BDs (e.g. Day-Jones et al.
A. C. Day-Jones et al.
2008; Burningham et al. 2009; Faherty et al. 2009; Zhang et al.
2010). In particular, white dwarf (WD) primaries can provide a
hard lower limit on the age of the system (from the WD cooling
age) and in the case of high-mass WDs (where the MS progenitor
star will have a short lifetime and the age will thus be essentially the
same as the cooling age of the WD), they could provide ages constrained at the 10 per cent level (Pinfield et al. 2006 and reference
Binary systems containing a WD and a BD however are
observationally rare (Farihi, Becklin & Zuckerman 2005), and
only a handful of such binaries have been identified. To
date, there are only five L dwarf + WD systems that have
been spectroscopically confirmed: GD 165B (L4; Zuckerman
& Becklin 1992), GD 1400 (L6/7; Farihi & Christopher 2004;
Dobbie et al. 2005), WD 0137−349 (L8; Maxted et al. 2006;
Burleigh et al. 2006), PG 1234+482 (L0; Steele et al. 2007; Mullally et al. 2007) and PHL 5038B (L8; Steele et al. 2009), previously
the latest type BD companion to a WD. There have also been several
BD companions to WDs found as cataclysmic variables (CVs, e.g.
Littlefair et al. 2008), although the nature of such objects means
that they may be less useful in the context of studies of typical BD
We present here results from our ongoing search to identify
widely separated BD companions to WDs, expanding on our earlier searches of the 2MASS and SuperCOSMOS (Day-Jones et al.
2008), to include the first results from our combined search of the
UKIDSS and SuperCOSMOS. We present here the discovery of
the first T dwarf + WD binary system, which we confirm through
common proper motion and spectroscopy.
In Section 2, we describe the ongoing search to identify and
spectroscopically confirm late T dwarfs in the UKIDSS Large Area
Survey (LAS). In Section 3, we describe our proper motion measurements of these confirmed T dwarfs and our techniques to search
for binary companions to these objects. In Section 4, we describe
the spectroscopy of a WD candidate companion and the resulting
constraints on its properties. In Section 5, we statistically assess
the likelihood that our new T dwarf + WD binary system is spurious. Section 6 determines constraints for the atmospheric properties
of the T dwarf, taking advantage of its benchmark age constraints
(from the WD primary). We also perform some basic spectral model
fits to the T dwarf spectrum and compare the resulting predictions.
Finally, in Section 7, we conclude with further discussion of the
system as a useful benchmark and comment on the direction of
future work.
2.1 Follow-up photometry and spectroscopy
Candidates are followed up in general with additional imaging in
the NIR and/or optical. This allows confirmation of the expected Tlike colours and rules out various forms of contamination (e.g. faint
M dwarfs with low signal-to-noise ratio (S/N) and blueward scattered NIR colours due to large photometric uncertainty, as well
as Solar system objects that can appear as non-detections due to
their motion). This follow-up has been performed using a variety of
telescope/instruments, including the Wide Field Camera (WFCAM)
and Fast Track Imager (UFTI) on the UKIRT and the Long-slit Intermediate Resolution Infrared Spectrograph (LIRIS) on the William
Herschel Telescope (WHT) (all NIR), as well as the ESO MultiMode Instrument (EMMI) and the ESO Faint Object Spectrograph
and Camera (EFOSC2) on the New Technology Telescope (optical). Spectroscopic confirmation of the LAS candidates was also
achieved using a number of facilities, including the Near Infrared
Camera and Spectrograph (NIRI) and the Gemini Near
Infrared Spectrograph (GNIRS) on the Gemini telescopes, and the
Infrared Camera Spectrograph (IRCS) on Subaru. The Near Infrared
Camera Spectrograph (NICS) on the Telescopio Nazionale Galileo
(TNG) and the UKIRT imager Spectrograph (UIST) were also used
for the brighter T dwarfs. Details of the follow-up imaging and spectroscopic strategies, as well as the relevant reduction and calibration techniques used, are further discussed in Lodieu et al. (2007),
Pinfield et al. (2008) and Burningham et al. (2010). Corresponding
spectral types were derived using the unified T dwarf classification
scheme of Burgasser et al. (2006) with an extension from Burningham et al. (2008) for the latest types.
2.2 T dwarf distances
T dwarf distances have been estimated using the absolute
magnitude–spectral type relations from Liu et al. (2006), assuming that the T dwarfs are single objects. We calculated absolute
magnitude (MJ ), choosing the J-band magnitude over the H and
K bands, as models have suggested that the J band may be less
sensitive to variations in metallicity and gravity than the H and K
bands (e.g. Marley et al. 2002). The uncertainties in the distance
were obtained by taking into account the error in the spectral type
(typically ±0.5) and the residuals of the polynomial fits from Liu
et al. (2006). Fig. 1 shows the spectral type distance distribution
for the spectroscopically confirmed sample of 49 T dwarfs that we
The UKIDSS LAS has been searched for T dwarfs using selection techniques based on the observed UKIDSS+SDSS colours of
previously identified T dwarfs, as well as theoretical predictions
for the cooler T eff = 400–700 K regime. The selection techniques
used to identify these T dwarfs are described in detail in Pinfield
et al. (2008). In this paper, we consider T dwarfs that were spectroscopically confirmed (Burningham, private communication) by
the Summer of 2008 (see Pinfield et al. 2008 and a subsample of
Burningham et al. 2010) and the sky coverage appropriate to this
sample includes the full LAS second data release, 72 per cent of
the new sky in the third data release and 66 per cent of the new
sky in the fourth data release. In total, this sample was identified in
890 deg2 of the LAS sky.
Figure 1. T dwarf distance versus spectral type for our spectroscopically
confirmed T dwarfs with second epoch imaging.
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation Discovery of a T dwarf + white dwarf binary
Figure 2. A vector point diagram showing the proper motions of our
spectroscopically confirmed T dwarfs with good proper motion estimates
(see Section 3). Shown are the different spectral types T4–6, T7–8 and
T8+ as squares, triangles and circles, respectively. ULAS J1459+0857 and
Wolf 940B are also highlighted.
consider in this work. It can be seen that the sample spans the spectral type and distance range of T2–9 and 12–80 pc, respectively.
This is comprised of two T2–3 dwarfs, 29 in the T4–6 range, 16
T6–8 dwarfs and two T8+ dwarfs. Due to the cooler temperatures
and thus fainter nature of later T dwarfs, it can be seen that we are
more sensitive to earlier T dwarfs out to farther distances.
The photometric follow-up programme provided second epoch
imaging for the LAS T dwarfs, which we combined with the LAS
images to give two epochs to calculate proper motion. We used the
IRAF routines GEOMAP and GEOXYTRAN to derive a geometric transformation between the two epoch images and to apply these transforms
to the T dwarf positions, respectively. Centroiding uncertainties
were calculated based on simulated data with appropriate Poisson
noise injected. The availability and quality of the measured proper
motions of the T dwarfs are dependent on several factors, including
the baseline between the epochs, the number of stars that can be used
for positional reference in each of the images, the S/N of both the T
dwarf and the reference stars, and the proximity of the T dwarfs to
the edge of the WFCAM detector array (in the first epoch images).
We were only able to measure the proper motions of 19 T dwarfs
from the 49 strong sample of spectroscopic confirmations that we
consider, using an average of 12 reference stars across a baseline of
0.5–1.5 Gyr. A vector-point diagram of the T dwarf proper motions
is shown in Fig. 2, where two T dwarfs found to have common
proper motion companions (see Section 3.1) are highlighted.
3.1 A search for common proper motion companions
We searched for candidate common proper motion companions to
our sample of T dwarfs with reliable proper motions, by selecting a
magnitude-limited sample of SuperCOSMOS sources, where R <
21 (the magnitude limit at which proper motions are measured) and
have accurately measured proper motions, such that PM/σ PM ≥
3. We searched around each of the 19 T dwarfs out to an angular
separation corresponding to 20 000 au, at the estimated minimum
distance of each T dwarf. We choose a separation limit of 20 000 au
in order to be sensitive to the detection of both widely separated
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation 707
MS and WD companions. It is fairly common to find BD + MS star
binaries with separations of ∼5000 au (Gizis et al. 2001; Pinfield
et al. 2006). However, a WD companion could have even wider
separations, when one considers any outward migration that would
have occurred during the post-MS mass-loss phase. We suggest that
the outward migration would likely be a factor of ∼4, as the initial
and final separation of a low-mass binary companion is directly
related to the change in mass of the host star, that is, M MS /M WD
(Jeans 1924; Zuckerman & Becklin 1987). To illustrate this, we
consider a WD of mass ∼0.65 M (roughly the mean of the WD
mass distribution). The progenitor mass would be ∼2.7 M (from
the initial–final mass relations of Dobbie et al. 2006; Catalán et al.
2008; Kalirai et al. 2008), such that M MS /M WD ∼ 4. Thus, for BD
+ MS binaries separated by 5000 au, we could expect their final
separation to be up to ∼20 000 au.
For each common proper motion companion candidate, we assumed a distance that was the same as the estimated distance of
the relevant T dwarf, and thus placed companion candidates on an
absolute B magnitude (MB ) versus B − R colour–magnitude diagram. Their positions were compared with those expected for MS
stars and WDs, following the approaches described in Clarke et al.
(2009) and Day-Jones et al. (2008), respectively. We then compared
their SDSS colours with respect to stellar populations, including
MS stars (Hipparcos; Perryman et al. 1997), M dwarfs (West et al.
2004), K subdwarfs (Yong & Lambert 2003) and WDs (McCook
& Sion 1999; Eisenstein et al. 2006). We found that only one MS
companion candidate was identified, Wolf 940, which had previously been identified serendipitously by Burningham et al. (2009)
and will not be discussed further. We also identified five candidate
WD companions to the T dwarfs in our sample that are common
proper motion (to within the uncertainties) and are consistent with
the WD sequence (see Figs 3 and 4), if assumed to be at the same
distance as their T dwarf companions.
The brightest of our WD candidates appears in the LSPM catalogue (LSPM J1459+0851, Lépine & Shara 2005) as a high
proper motion object, although it has not been previously studied spectroscopically. The T dwarf associated with this object is
ULAS J1459+0857, which has been spectroscopically typed as a
T4.5 ± 0.5 dwarf (Burningham et al. 2010). This pair is highlighted
in Figs 3 and 4.
3.2 Possible contamination
To assess the possibility that our candidate WD companions may
be dominated by high-velocity background objects, such as metalpoor, halo K subdwarfs (which could populate the same colour and
proper motion space that we search in this work), we have derived
space motion estimates, assuming that our candidates are K subdwarfs. We used the relations of Ivezić et al. (2008) to estimate an
absolute r magnitude (Mr ) from g − i colour, for a metallicity
range (for subdwarfs) of −0.5 and −2.5, and thus obtained distance
constraints applicable, if these objects are subdwarfs. Assuming, for
simplicity, that they have a zero radial velocity, we then estimated
space motions to assess potential halo membership. All except one
candidate would have space motions of 2000–12 000 km s−1 , which
are thus not consistent with a Galactic halo population. The exception is close enough, such that its kinematics are consistent
with a background subdwarf member. Indeed, this is one of the
widest separated candidates (400 arcsec), and the level of background contamination for such separation (and volume) approaches
1, even for the low-density halo luminosity function (Gould 2003).
A. C. Day-Jones et al.
Figure 3. A colour–magnitude diagram of WDs from McCook & Sion (1999) with known parallax (crosses). Photometry is on the SuperCOSMOS system.
Overplotted are model cooling tracks (see main text) for WD masses of 0.5, 0.7 and 1.2 M (dotted, dashed and dot–dashed lines, respectively). Also
overplotted is our WD selection region (two solid lines), along with our candidate WD companions (large stars). LSPM J1459+0851 is circled for reference.
Figure 4. A two-colour diagram in SDSS colours showing populations of MS stars (blue points), M dwarfs (orange triangles), WDs (green asterisks) and K
subgiants (red squares), with our WD candidates overplotted as black stars. LSPM J1459+0851 is circled for reference.
However, despite one possible halo contaminant, the majority of the
candidates cannot be explained by such contamination.
A more likely source of error potentially leading to misidentification amongst the candidates is a proper motion uncertainty, since the
ratio of proper motion to proper motion uncertainty for some of the
candidates is in the 3–5 range, and one would thus expect that some
fraction of the sample have proper motions that have scattered to
large values. Given this likely source of contamination, we choose
not to present details of all five of our companion candidates at this
stage. We prefer to first establish their nature through spectroscopic
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation Discovery of a T dwarf + white dwarf binary
study, and in that way, confirm them (or otherwise) as genuine WD
companions. We have so far only obtained good spectroscopy for
both components of the binary containing LSPM J1459+0851 and
thus focus the remainder of the paper on this system.
J 1 4 5 9 +0 8 5 1
Spectroscopic observations of LSPM J1459+0851 were obtained
with FORS2 on the Very Large Telescope on 2009 May 15 and
21, with Directors Discretionary Time in program 282.C-5069(A).
We used the long-slit mode in the optical wavelength range 3300–
8000 Å with a dispersion of 50 and 55 Å mm−1 , respectively, for
the ranges 3300–6210 Å (corresponding to the B grism) and 5120–
8450 Å (corresponding to the RI grism), giving a resolution of R ∼
1200. Three integrations of 600 s were taken, giving a total exposure
time of 30 min for the B grism and two integrations of 360s, totaling
to 12 min in the RI grism. Sky flats, arc frames and the spectra of
a featureless (DC) WD as well as a standard F-type star and a
hydrogen-rich (DA) WD were taken during the same night at a
similar airmass to the target so as to provide wavelength, flux and
telluric calibrations.
Standard IRAF routines were used to reduce the spectra, including
flat-fielding and cosmic-ray removal. The spectra were extracted
with APALL, using a Chebyshev function to fit the background and a
third-order Legendre function to trace the fit to the spectrum. The
wavelength calibration was done using the spectrum from HgCdHeAr and HgCdHeNeAr arc lamps for the B and RI grisms, respectively, and using IDENTIFY to reference the arc lines, along with
the DISPCOR routine to correct the dispersion of the spectrum. The
resulting spectra of both LSPM J1459+0851 and the standard were
divided by the smooth spectrum of the DC WD, which has no intrinsic spectral features, enabling correction for the instrumental
response. The standard stars (one for each grism) were then used
to flux calibrate the spectrum. The two spectra were then stitched
together in the overlapping sections and normalized at 6000 Å. The
final spectrum of LSPM J1459+0851 is shown in Fig. 5.
It can be seen that some residual tellurics remain and are highlighted for reference. They do not, however, affect the subsequent
analysis in any way, since they do not overlap with features used
directly to assess WD properties. The general spectral shape is quite
blackbody-like, consistent with a WD or perhaps a very metal poor
Figure 6. Optical spectrum in the region of Hα for the WD LSPM
J1459+0851. For comparison, the spectra of three similar cool, hydrogenrich, magnetic WDs, WD 0011−134 (B ∼ 16.7 MG), WD 1330+015 (B ∼
7.4 MG) and WD 0503−174 (B ∼ 7.3 MG) are also shown, with magnetic
field strength decreasing from top to bottom.
subdwarf (e.g. Jao et al. 2008). However, the overall strength of
the Hα line and the peak of the blackbody-like continuum are only
consistent with a relatively cool example of a WD (Kilic et al.
2006). We also compare LSPM J1459+0851 to the spectra of three
other very cool, DA WDs, WD 0011−134, WD 1330+015 and
WD 0503−174, taken from Bergeron, Ruiz & Leggett (1992) and
Bergeron, Ruiz & Leggett (1993). They have corresponding values
of T eff of 6000 ± 150, 7450 ± 200 and 5230 ± 140 K, respectively,
and are shown in Fig. 6. Although the spectra of LSPM J1459+0851
are noisier, it can be seen that the extent of the Hα feature is consistent with a cool, DA WD. We further assess its properties in more
detail in the following sections.
4.1 Synthetic photometry and fitting procedure
Figure 5. The optical spectra of LSPM J1459+0851, normalized at 6000 Å.
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation As previously noted, the spectrum shows a lack of strong H lines,
which would be expected for a hotter more typical WD. Photometric fitting of the full optical–NIR spectral energy distribution
(SED) is thus optimal for determining the effective temperature
of the WD, using model fits to the SDSS+UKIDSS photometry
and assuming a distance equal to that of the T dwarf companion.
The spectra are consistent with a cool WD, showing no strong features in the spectrum blueward of 6000 Å and just a hint of Hα.
We performed a fit to the available photometry of the WD (from
SDSS, USNO and UKIDSS) using the atmospheric model codes of
Bergeron et al., which are described at length in Bergeron, Wesemael & Beauchamp (1995, with updates given in Bergeron, Leggett
& Ruiz 2001; Bergeron et al. 2005). These models are in local
A. C. Day-Jones et al.
thermodynamic equilibrium and allow energy transport by convection and can be calculated with arbitrary amounts of hydrogen and
helium. Synthetic colours1 were obtained using the procedure outlined in Holberg & Bergeron (2006) based on the Vega fluxes taken
from Bohlin & Gilliland (2004).
The method used to fit the photometric data is similar to that described in Bergeron et al. (2001), which we briefly summarize here.
We first transform the magnitudes in each bandpass into observed
average fluxes f mλ using the following equation:
m = −2.5 log fλm + cm ,
fλ Sm (λ)dλ
= 0 ∞
Sm (λ)dλ
The transmission functions Sm (λ) along with the constants cm for
each bandpass are described in Holberg & Bergeron (2006) and
references therein. To make use of all the photometric measurements simultaneously, we convert the magnitudes into observed
fluxes using equation (1) and compare the resulting energy distributions with those predicted from our model atmosphere calculations.
Thus, we obtain a set of average fluxes f mλ , which can now be compared with the model fluxes. These model fluxes are also averaged
over the filter bandpasses by substituting f λ in equation (2) for the
monochromatic Eddington flux H λ . The average observed fluxes f mλ
and model fluxes H mλ , which depend on T eff , log g and N(He)/N(H),
are related by the equation
Hλm ,
fλm = 4π
where R/D is the ratio of the radius of the star to its distance from
Earth. Our fitting procedure relies on the non-linear least-squares
method of Levenberg–Marquardt, which is based on a steepest descent method. The value of χ 2 is taken as the sum over all bandpasses of the difference between both sides of equation (3), properly
weighted by the corresponding observational uncertainties. In our
fitting procedure, we consider only T eff and the solid angle free
parameters. As discussed in Bergeron et al. (2001), the energy distributions are not sensitive enough to surface gravity to constrain the
value of log g, and thus for WDs with no parallax measurement, as
is the case here, we simply assume log g = 8.0, which is consistent
with the distance estimate of the T dwarf companion. Our best fit
for a pure-H atmosphere arises from a T eff = 5535 ± 45 K, and is
shown in Fig. 7.
Figure 7. Top panel: Fit of the energy distribution with pure-H models. The
observed ugriz and JHK fluxes are represented by error bars, while the model
fluxes averaged over the filter bandpasses are indicated by the filled circles.
The model monochromatic fluxes are shown by a solid line. Bottom panel:
Normalized spectrum near Hα with the synthetic profiles interpolated at the
parameters obtained from the energy distribution fits, assuming a pure-H
atmospheric composition with a model fit at T eff = 5535 ± 45 K and log g
= 8.0 with an offset dipole magnetic field computed with the parameters
indicated in the figure.
polarization of the radiation is neglected as we are mainly interested
in the total monochromatic intensity. The effect of the magnetic
field on the continuum opacity is also neglected. The emergent
spectrum is then
obtained from an integration over the surface of
the star (Hν ∝ Iν μdμ) for a particular geometry of the magnetic
field distribution. We note that in this procedure, limb darkening is
explicitly taken into account because of the integration over μ.
We use the offset dipole model to model the magnetic field of
the WD (Achilleos & Wickramasinghe 1989). In this model, the
magnetic field is generated by a dipole. At the surface of the star (of
radius unity), the strength of the magnetic field is simply given by
4.2 A magnetic spectrum
The shape and weakness of the Hα line gives a poor fit to a basic
5535 K model spectra, and we thus investigated the possibility that
the WD could be magnetic with the Hα line affected by Zeeman
splitting. The line opacity was calculated as the sum of the individual
resonance-broadened Zeeman components. The line displacements
and strengths of the Zeeman components of Hα are taken from the
tables of Kemic (1974), and the total line opacity is normalized to
that resulting from the zero-field solution. The specific intensities
at the surface, I(ν, μ, τ ν = 0), are obtained by solving the radiative
transfer equation for various field strengths and values of μ(μ =
cos θ, where θ is the angle between the angle of propagation of
light and the normal to the surface of the star). In doing so, the
B = 0.5Bd (3 cos2 θ + 1)1/2 ,
where Bd is the dipole field strength and θ is the standard angle in
spherical coordinates (θ = 0 at the pole). For a given value of Bd ,
the flux received at the Earth will also depend on the viewing angle
i between the dipole axis and the line of sight (i = 0 for a pole-on
view). However, in this particular model, the dipole is also offset
from the centre of the star in an arbitrary direction. To simplify the
calculation, we assume that the dipole is offset parallel to the dipole
axis. In this case, the value of the offset is measured from the centre
The synthetic colours can be obtained at http://www.astro.umontreal.ca/
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation Discovery of a T dwarf + white dwarf binary
of the star and is denoted by az (in units of stellar radius). Note that
with the offset dipole models, the value of the dipole field strength,
Bd , is no longer equal to the value of the polar field strength.
We computed a series of synthetic spectra based on a model
atmosphere of pure hydrogen with T eff = 5535 K and log g = 8.0,
while varying Bd , i and az . We display in Fig. 7 the model which
best reproduced the observed line profile with Bd = 2.0 MG, i =
45◦ and az = −0.20. It should be noted that varying the inclination
angle i produced only slight variations in the line profiles and as
stated in Bergeron et al. (1992), it is not possible to constrain i from
observed line profiles alone. The Zeeman-split Hα line in LSPM
J1459+0851 can be compared to the same feature in other cool
WDs of higher magnetic field strengths in Fig. 6.
4.3 White dwarf age and progenitor lifetime
We start by considering the simple case of a non-magnetic WD
for a T eff and log g of 5535 K and 8.0, respectively, for which
we calculate a mass of 0.585 M (Fontaine, Brassard & Bergeron
2001). The corresponding cooling age of a 0.585-M WD, with a
T eff = 5535 K, was then calculated as 3 Gyr, using the isochrones
of Fontaine et al. (2001). The total age of the WD comprises
both the cooling time and its progenitor lifetime on the MS. In
order to constrain the MS lifetime, we accessed the initial–final
mass relations of Ferrario et al. (2005), Catalán et al. (2008) and
Kalirai et al. (2008) to estimate a likely, initial-mass constraint for
the MS progenitor star of 1.50–1.75 M . We then used the tracks of
Lachaume et al. (1999) to estimate the MS lifetime for stars of such
mass as 1.8–3.0 Gyr. It should be noted, however, that the model
tracks converge for masses <2 M , and as a result the ages of these
objects can be largely uncertain (up to 10 Gyr).
The strong magnetic field present also provides an additional
factor to consider when assessing age. The origin of such strong
WD magnetic fields is not fully understood but is thought to have
arisen in one of the following two favoured scenarios:
(i) From a single star. The magnetic field is thought to derive
from a massive, magnetic progenitor of ∼1.5–8 M , typically an
Ap or Bp star. The magnetic field is then maintained through the MS
evolution to the WD phase by flux conservation (Wickramasinghe
& Ferrario 2000).
(ii) From the merger of two stellar cores in a common envelope
(CE) event or the merger of two degenerate objects. During the CE
phase the orbits of the two cores spiral in closer together through
frictional forces causing differential rotation, which coupled with
convection in the cores, creates a stellar-magnetic dynamo (Tout
& Pingle 1992). Close, but separated, cores form CVs, but some
cores coalesce and cool to form a magnetic WD. It may also be
possible that two very close WDs emerge from the CE phase, such
as G62–46 (Bergeron et al. 1993), where one component is highly
In general, it is observed that magnetic WDs have larger masses
than the more typical non-magnetic WDs (Liebert 1988). There
are two possible hypotheses to explain this. First, in accordance
with the favoured scenario for the formation of isolated magnetic
WDs (Wickramasinghe & Ferrario 2005), the progenitor was more
massive, leading to a massive WD. In this case, the magnetic field
has no effect during the progenitor evolution. Secondly, the effect
of the magnetic field has an impact on the stellar evolution, such
that it could inhibit mass loss (Wickramasinghe & Ferrario 2000),
leading to a more massive core and a longer progenitor lifetime. If
we consider the possibility that our WD could be of higher mass, for
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation 711
Table 1. Parameters of the WD LSPM J1459+0851.
μ RA
μ Dec.
T eff
log g
WD age
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
14h 59m 32.s 05
+08◦ 51 28. 1
20.74 ± 0.08
19.50 ± 0.01
18.99 ± 0.01
18.76 ± 0.05
18.71 ± 0.03
18.14 ± 0.02
17.90 ± 0.02
17.65 ± 0.04
17.79 ± 0.06
−170 ± 3 mas yr−1a
−42 ± 6 mas yr−1a
5535 ± 45 K
8.0 dex
0.585 M
>4.8 Gyr
a USNO-B1.
example, 0.8 M (the mean mass of a highly magnetic WD; Kawka
et al. 2007), then the cooling time would be longer, around 6 Gyr
(Fontaine et al. 2001). In this case, the progenitor (if a single star)
would be around 3.5 M (Catalán et al. 2008) and the progenitor
MS lifetime would be 0.3 Gyr (assuming normal models). However,
there is also evidence that many magnetic WDs have masses closer
to the peak of the non-magnetic mass distribution (Tout et al. 2008).
As we cannot know which scenario is responsible for the observed
magnetic field, nor can we measure more accurately the mass of
the WD, we do not know the effects this may have had on the
MS evolution. In any case, both scenarios for a magnetic and a
non-magnetic WD result in ages greater than 4.8 Gyr (the cooling
age plus the minimum progenitor lifetime of a non-magnetic WD)
and we thus choose to adopt this as the minimum age for LSPM
J1459+0851. Properties of the new WD are given in Table 1.
5 L S P M J 1 4 5 9 +0 8 5 1 – U L A S J 1 4 5 9 +0 8 5 7 :
In order to determine if this new system is a bona fide binary system,
we have statistically assessed the likelihood that two such objects
could be a line-of-sight association with photometry and proper motion consistent with binarity by random chance. We first calculated
the total region of sky around our T dwarf corresponding to the coverage encompassed by the projected line-of-sight separation of the
WD. We then combined this with the T dwarf distance constraint
(43–69 pc), allowing for the possibility that the T dwarf might itself be an unresolved binary (at a greater distance), to estimate a
volume of sky in which WDs might be line-of-sight contamination.
We then used the number density of WDs (e.g. Schröder, Pauli &
Napiwotzki 2004) to estimate that we would expect only 0.003 68
WDs in this volume of space.
To factor in the probability that two objects might have a common
proper motion at the level of our measurements, we downloaded a
magnitude-limited sample (R < 21, the same as our initial selection;
A. C. Day-Jones et al.
∼4100 au and more akin to the more common type of BD + MS
binaries (<5000 au; Pinfield et al. 2008).
Table 2. Parameters of the binary system.
Separation on sky
Estimated distance
Estimated line of sight separation
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
a Assuming
385 arcsec
43–69 pca
16 500–26 500 aua
6 P RO P E RT I E S O F U L A S J 1 4 5 9 +0 8 5 7
In order to estimate the T eff of ULAS J1459+0857, we used spectral
type–T eff determinations from table 6 of Golimowski et al. (2004)
as a guide. These determinations are for BDs with known parallaxes
(e.g. Tinney 1996; Dahn et al. 2002; Leggett et al. 2002; Tinney,
Burgasser & Kirkpatrick 2003; Knapp et al. 2004; Vrba et al. 2004)
and SED constraints over a broad wavelength range, for which
reliable bolometric flux measurements and luminosities are thus
available. The main source of uncertainty in these T eff values comes
from a lack of strong age constraints and the resulting evolutionary
model radius uncertainties (up to ∼30 per cent for ages >0.1 Gyr;
Burrows et al. 1997; Baraffe et al. 1998; Chabrier at al. 2000). By
considering the variety of T eff ranges calculated for the T4.5 ± 0.5
dwarfs, which have assumed a range of possible ages from 0.1 to
10 Gyr (± ∼300 K), we estimate that for an age range of 4–10 Gyr,
ULAS J1459+0857 has T eff in the range 1200–1500 K.
We also used the Lyon Group COND models (Baraffe et al. 2003)
to estimate the physical properties of ULAS J1459+0857, allowing for the possibility that it could be a single object or itself an
unresolved binary, with a distance in the range 43–69 pc. We calculated absolute magnitudes for ULAS J1459+0857 based on this
distance range for ages 4–10 Gyr. We then obtained mass and log g
estimates in the range 0.064–0.075 M and 5.4–5.5 dex, respectively, assuming solar metallicity, by using a linear interpolation
between the model grid points. We also consider the evolutionary
models of Burrows et al. (1997, 2001) and Burrows, Sudarsky &
Hubeny (2006) to estimate a mass of 0.064–0.075 M , if the T
dwarf is actually metal poor ([M/H] − 0.5 dex), which is similar to
that of the solar metallicity COND models. Both models also indicate a high gravity (log g = 5.5) for the observable J − K colours
and temperature range of ULAS J1459+0857. We also compare
the optical + NIR colours of ULAS J1459+0857 in comparison
with other spectroscopically confirmed T dwarfs from the UKIDSS
LAS (Burningham et al. 2010). Fig. 9 shows this complement of T
dwarfs compared with the spectral type. Whilst the NIR colours of
ULAS J1459+085, in general look fairly typical for a T4.5 dwarf,
that the T dwarf is a singular or unresolved binary.
see Section 3.1) of objects from the SuperCOSMOS Science
Archive. We applied a limit to the proper motion uncertainty of
<50 mas yr−1 and required objects to lie in the colour range 1 <
B − R < 3 (where we expect contaminant, MS stars to populate).
This sample of 10 360 sources was selected from within one degree
of the T dwarf as to provide a representative sample of objects in
the area of sky in which we find our binary system. We then placed
these objects on a colour–magnitude diagram and selected only objects that occupied a region populated by MS stars, when placed at
the distance range estimated for the T dwarf. Of the 140 objects that
were selected in this way, 13 had proper motion consistent (at the
1σ level) with the T dwarf, leading to a probability of 0.092 that
a contaminant star could have proper motion common with the T
The chance of finding a WD with the same line-of-sight separation of LSPM J1459+851 from ULAS J1459+0857, where both
components share a common proper motion and are consistent with
lying at the same distance as that estimated for the T dwarf, is thus
0.003 68 × 0.092 = 0.0003. We thus conclude that these objects
form a genuine binary system. Since we searched for companions
to a total of 19 T dwarfs, we estimate that the overall chance of finding a spurious system in our sample is 0.0064 and that the systems
identified are likely real binaries. Properties of the new WD + T
dwarf system are given in Table 2 and a finding chart is presented
in Fig. 8. The wide separation of the system (16 500 au, assuming
it to be a singular object and not an unresolved binary) is similar
to the widest known BD + MS binary systems (e.g. Faherty et al.
2009; Zhang et al. 2010), although prior to the post-MS mass-loss
phase of the primary, the separation would have been substantially
less. Indeed, we expect that the initial separation of the two components was a factor of ∼4 closer (see Section 3.1) in the region of
Figure 8. A UKIDSS J-band finder chart showing the position of the T dwarf (square) and the WD (circle), the scale of the image is 8 × 3.5 arcmin2 .
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation Discovery of a T dwarf + white dwarf binary
Figure 9. Colour against T dwarf spectral type for UKIDSS T dwarfs (Burningham et al. 2010). ULAS J1459+0857 is shown as a large, filled red circle.
there is some evidence of relative z band enhancement, which could
be the result of the older, higher gravity nature of this object (Pinfield
et al. 2008).
6.1 Initial model testing with benchmark observations
To provide a first-pass test of model atmosphere predictions, we
used the dust-free COND models of Baraffe et al. (2003) to provide
theoretically informed best-guess constraints of the physical properties of the T dwarf in our binary. We made comparisons between
the observed T dwarf spectrum and synthesized spectroscopy for
two values of T eff (1200 and 1500 K), three values of log g (5.0,
5.25 and 5.5) and three metallicities (+0.3, 0.0 and −0.5, representing the range observed in the Galactic disc; e.g. Valenti &
Fischer 2005; Gray et al. 2006; Holmberg, Nordstróm & Andersen 2007; Jenkins et al. 2008). The model spectra were generated
using the atmospheric radiative transfer code, Phoenix (which is described in detail in Hauschildt, Allard & Baron 1999; Allard et al.
2001), which includes the updated water molecular opacity list from
Barber et al. (2007) and new solar abundances from Asplund et al.
(2009). The model includes the effect of condensation in the chemical equilibrium but ignores the effects of dust opacities. Theoretical
spectra were normalized to the observations in the peak of the J
band, and the resulting comparisons are shown in Figs 10 and 11.
The long red-dotted, short-dashed green and the long-dashed blue
lines represent log g = 5.0, 5.25 and 5.5, respectively.
Table 3 gives a summary of the reduced chi-squared (χ 2ν ) values
calculated for these comparisons. These values range from ∼1.75
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation to 3.5, and we note that none of them are close to 1.0. This is to
be expected, since there are known shortcomings in the models that
introduce differences significantly greater than the measurement
uncertainties – e.g. incomplete methane opacities and a poor understanding of the observed J-band brightnening around T3 (e.g. Knapp
et al. 2004). As such, we only use these χ 2ν values as an additional
quality indicator for the overall model-observation comparison.
Our χ 2ν analysis reveals that the best-fitting atmospheric model
has an effective temperature of 1200 K, a subsolar metallicity and
a high gravity of log g = 5.5. However, comparison by eye shows
clearly that this fit across the whole spectral range is not particularly good, especially around the peak flux regions of the H and
K bands. This best χ 2ν value presumably comes from the better
fit to the J band, compared to the other models; however, a better fit (by eye) to the overall profile of the spectra comes from a
1500-K, solar metallicity, high-gravity model. This suggests that
ULAS J1459+0857 probably has solar to slightly subsolar metallicity and high gravity. This is instructive at least that the observed
properties of ULAS J1459+0857 appear to be in general agreement
with those predicted by evolutionary models. While the models
clearly have some shortcomings, they appear to be making progress,
such that model properties for a mid-T dwarf of age >4 Gyr seem
to be broadly consistent with benchmark properties.
7 S U M M A RY
This is the first discovery of a T dwarf + WD binary system and
an example of an evolved, high-gravity BD. During the MS phase
A. C. Day-Jones et al.
Figure 10. Spectral model comparisons to ULAS J1459+0857 (black solid line) for T eff = 1200 K and [M/H] = +0.3, 0.0 and −0.5, from the top to bottom
panel, with log g = 5.0, 5.25 and 5.5 shown as a long red dotted line, short-dashed green line and the long-dashed blue line, respectively.
Figure 11. Spectral model comparisons to ULAS J1459+0857 (black solid line) for T eff = 1500 K and [M/H] = +0.3, 0.0 and −0.5, from the top to bottom
panel, with log g = 5.0, 5.25 and 5.5 shown as a long red dotted line, short dashed green and the long dashed blue lines, respectively.
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation Discovery of a T dwarf + white dwarf binary
Table 3. χ 2ν fits to model spectra.
log g
T eff = 1500 K
T eff = 1200 K
Table 4. Parameters of ULAS J1459+0857.
Further observations that would have clear benefits for this system
include a parallax measurement of either component to yield an
accurate distance and fuller spectral coverage of the T dwarf to
constrain its mid-IR and optical spectral morphology. In general,
and particularly in the optical z band, it would be desirable to
assess in more detail spectroscopic features and trends that may
be sensitive to higher surface gravity and age. A parallax distance
combined with accurate knowledge of the T dwarf bolometric flux
would offer significant improvements on the T eff constraints for this
object (e.g. Burningham et al. 2009).
An accurate distance would also facilitate greatly improved radius
and mass constraints for the WD, and thus a better constraint on the
cooling age and progenitor lifetime. More detailed studies of both
binary components are clearly crucial to maximize the effectiveness
of the benchmark BD component.
μ RA
μ Dec.
Spectral type
T eff
log g
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . ..
. . .. . .. . .. . ..
. . .. . .. . .. . ..
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
. . .. . .. . .. . .
14h 59m 35.s 25
+08◦ 57 51. 20
43–69 pca
21.17 ± 0.15
19.24 ± 0.06
17.93 ± 0.02
17.94 ± 0.05
17.92 ± 0.08
17.98 ± 0.04
17.93 ± 0.04
18.04 ± 0.03
−149 ± 33 mas yr−1
−45 ± 33 mas yr−1
T4.5 ± 0.5
0.060–0.072 Mb
0.064–0.075 Mc
1200–1500 K
5.4–5.5 dex
a Assuming
that the T dwarf is a singular or unresolved binary.
the Lyon group COND evolutionary models.
c From the Burrows evolutionary models.
b From
of the primary, its separation would have been similar to the bulk
population of BD+MS binaries, although the separation must have
grown significantly (to its current value) during the post-MS massloss phase of the primary.
The WD provides vital age constraints for the binary system
from its cooling age combined with a minimum estimate of its MS
progenitor lifetime, and the resulting minimum age of 4.8 Gyr for
the T dwarf allows a robust constraint on its surface gravity of
log g = 5.4–5.5. As such ULAS J1459+0857 is a representative
old, high-gravity benchmark BD. A comparison with the bulk popualtion of UKIDSS T dwarfs shows some indication that z band flux
enhancement may be an observational characteristic of older highgravity mid-T dwarfs, which could be a useful factor in attempts
to understand the formation history of substellar objects, and more
generally, this T dwarf can contribute to our understanding of substellar properties by providing a useful test-bed for theoretical model
This system is an example of how wide BD binary companions
to WDs make good benchmark objects, which will help test model
atmospheres and may provide independent means to calibrate BD
properties of field objects (Pinfield et al. 2006). This is the first of
our candidate systems and we expect to find many more as we mine
more the UKIDSS sky and with our ongoing efforts to search still
deeper by combining UKIDSS BDs with the SDSS for WDs.
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation ACD-J, JSJ and MTR would like to acknowledge the support of
the grant from CONICYT and the partial support from the Center
for Astrophysics FONDAP and Proyecto Basal PB06 (CATA). This
work was also suported in part by the NSERC Canada and by
the Fund FQRNT (Québec). PB is a Cottrell Scholar of Research
Corporation for Science Advancement. JG is supported by RoPACS,
a Marie Curie Initial Training Network funded by the European
Commission’s Seventh Framework Programme. SC is supported by
a Marie Curie Intra-European Fellowship within the 7th European
Community Framework Programme.
This publication has made use of the the data obtained from
the SuperCOSMOS Science Archive, prepared and hosted by the
Wide Field Astronomy Unit, Institute for Astronomy, University of
Edinburgh, which is funded by the STFC. We also acknowledge
the use of data from the Sloan digital archive, which is funded
by the Alfred P. Sloan Foundation, the Participating Institutions,
the National Aeronautics and Space Administration, the National
Science Foundation, the U.S. Department of Energy, the Japanese
Monbukagakusho, and the Max Planck Society. We have used data
from the Large Area Survey, including those from the data release
4 (Warren et al. 2007). We also acknowledge the use of DENIS
and the SIMBAD data base, operated at CDS, Strasbourg, France.
ACD-J is supported by a FONDECYT postdoctorado, under project
no. 3100098.
Achilleos N., Wickramasinghe D. T., 1989, ApJ, 346, 444
Allard F., Hauschildt P. H., Alexander D. R., Tamanai A., Schweitzer A.,
2001, ApJ, 556, 357
Asplund M., Grevesse N., Sauval A. J., Scott P., 2009, ARA&A, 47, 481
Baraffe I., Chabrier G., Allard F., Hauschildt P. H., 1998, A&A, 337, 403
Baraffe I., Chabrier G., Barman T. S., Allard F., Hauschildt P. H., 2003,
A&A, 702, 701
Barber R. J., Banerjee D. P. K., Ashok N. M., Tennyson J., 2007, in Corradi
R. L. M., Munari U., eds, ASP Conf. Ser. Vol. 363, The Nature of V838
Monocerotis and Its Light Echo. Astron. Soc. Pac., San Francisco, p. 95
Bergeron P., Ruiz M. T., Leggett S. K., 1992, ApJ, 400, 315
Bergeron P., Ruiz M. T., Leggett S. K., 1993, ApJ, 407, 733
Bergeron P., Wesemael F., Beauchamp A., 1995, PASP, 107, 1047
Bergeron P., Leggett S. K., Ruiz M. T., 2001, ApJS, 133, 413
Bergeron P., Ruiz M. T., Hamuy M., Leggett S. K., Currie M. J., Lajoie C.
P., Dufour P., 2005, ApJ, 625, 838
Bohlin R. C., Gilliland R. L., 2004, AJ, 127, 3508
Burgasser A. J., Geballe T. R., Leggett S. K., Kirkpatrick J. D., Golimowski
D. A., 2006, ApJ, 637, 1067
A. C. Day-Jones et al.
Knapp G. R. et al., 2004, AJ, 127, 3553
Lachaume R., Dominik C., Lanz T., Habing H. J., 1999, A&A, 348, 897
Leggett S. K. et al., 2002, ApJ, 564, 452
Leggett S. K. et al., 2009, ApJ, 695, 1517
Lépine S., Shara M. M., 2005, AJ, 129, 1483
Liebert J., 1988, PASP, 100, 1302
Littlefair S. P., Dhillon V. S., Marsh T. R., Gänsicke B. T., Southworth J.,
Baraffe I., Watson C. A., Copperwheat C., 2008, MNRAS, 388, 1582
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
Lodieu N. et al., 2007, MNRAS, 379, 1423
McCook G. P., Sion E. M., 1999, ApJS, 121
Marley M. S., Seager S., Saumon D., Lodders K., Ackerman A. S., Freedman
R. S., Fan X., 2002, ApJ, 568, 335
Maxted P. F. L., Napiwotzki R., Dobbie P. D., Burleigh M. R., 2006, Nat,
442, 543
Mullally F., Kilic M., Reach W. T., Kuchner M. J., von Hippel T., Burrows
A., Winget D. E., 2007, ApJS, 171, 206
Perryman M. A. C. et al., 1997, A&A, 323, 49
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,
Pinfield D. J. et al., 2008, MNRAS, 390, 304
Saumon D. et al., 2007, ApJ, 656, 1136
Schröder K. P., Pauli E. M., Napiwotzki R., 2004, MNRAS, 354, 727
Steele P. R., Burleigh M. R., Dobbie P. D., Barstow M. A., 2007, MNRAS,
382, 1804
Steele P. R., Burleigh M. R., Farihi J., Gänsicke B. T., Jameson R. F., Dobbie
P. D., Barstow M. A., 2009, A&A, 500, 1207
Tinney C. G., 1996, MNRAS, 281, 644
Tinney C. G., Burgasser A. J., Kirkpatrick J. D., 2003, AJ, 126, 975
Tout C. A., Pringle J. E., 1992, MNRAS, 259, 604
Tout C. A., Wickramasinghe D. T., Liebert J., Ferrario L., Pringle J. E.,
2008, MNRAS, 387, 897
Valenti J. A., Fischer D. A., 2005, ApJS, 159, 141
Vrba F. J. et al., AJ, 127, 2948
Warren S. J. et al., 2007, MNRAS, 381, 1400
West A. A. et al., 2004, AJ, 128, 426
Wickramasinghe D. T., Ferrario L., 2000, PASP, 112, 873
Wickramasinghe D. T., Ferrario L., 2005, MNRAS, 356, 1576
Yong D., Lambert D. L., 2003, PASP, 115, 22
Zhang Z. H. et al., 2010, MNRAS, 404, 1817
Zuckerman B., Becklin E. E., 1987, ApJ, 319, 66
Zuckerman B., Becklin E. E., 1992, ApJ, 386, 260
Burleigh M. R., Hogan E., Dobbie P. D., Napiwotzki R., Maxted P. F. L.,
2006, MNRAS, 373, L55
Burningham B. et al., 2008, MNRAS, 391, 320
Burningham B. et al., 2009, MNRAS, 395, 1237
Burningham B. et al., 2010, MNRAS, in press
Burrows A. et al., 1997, ApJ, 491, 856
Burrows A., Hubbard W. B., Lunine J. I., Liebert J., 2001, Rev. Mod. Phys.,
73, 719
Burrows A., Sudarsky D., Hubeny I., 2006, ApJ, 640, 1063
Catalán S., Isern J., Garcı́a-Berro E., Ribas I., 2008, MNRAS, 387, 1693
Chabrier G., Baraffe I., Allard F., Hauschildt P., 2000, ApJ, 542, 464
Clarke J. R. A. et al., 2010, MNRAS, 402, 575
Dahn C. C. et al., 2002, AJ, 124, 1170
Day-Jones A. C. et al., 2008, MNRAS, 388, 838
Delorme P. et al., 2008, A&A, 482, 961
Dobbie P. D., Burleigh M. R., Levan A. J., Barstow M. A., Napiwotzki R.,
Holberg J. B., Hubeny I., Howell S. B., 2005, MNRAS, 357, 1049
Dobbie P. D. et al., 2006, MNRAS, 369, 383
Eisenstein D. J. et al., 2006, ApJS, 167, 40
Faherty J. K., Burgasser A. J., West A. A., Bochanski J. J., Cruz K. L., Shara
M. M., Walter F. M., 2009, AJ, 137, 1
Farihi J., Christopher M., 2004, AJ, 128, 1868
Farihi J., Becklin E. E., Zuckerman B., 2005, ApJS, 161, 394
Ferrario L., Wickramasinghe D., Liebert J., Williams K. A., 2005, MNRAS,
361, 1131
Fontaine G., Brassard P., Bergeron P., 2001, PASP, 113, 409
Gizis J. E., Kirkpatrick J. D., Burgasser A., Reid I. N., Monet D. G., Liebert
J., Wilson J. C., 2001, ApJ, 551, 163
Golimowski D. A. et al., 2004, AJ, 127, 3516
Gould A., 2003, ApJ, 583, 765
Gray R. O. et al., 2006, AJ, 132, 161
Hauschildt P. H., Allard F., Baron E., 1999, ApJ, 512, 377
Holberg J., Bergeron P., 2006, AJ, 132, 1221
Holmberg J., Nordstróm B., Andersen J., 2007, A&A, 475, 519
Ivezić Ž., Sesar B., Jurić M., Bond N., Dalcanton J., Rockosi C. M., Yanny
B., 2008, ApJ, 684 287
Jao W. C., Henry T. J., Beaulieu T. D., Subasavage J. P., 2008, AJ, 136, 840
Jeans J. H., 1924, MNRAS, 85, 2
Jenkins J. S., Jones H. R. A., Pavlenko Y., Pinfield D. J., Barnes J. R.,
Lyubchik Y., 2008, A&A, 485, 571
Kalirai J. S., Hansen B. M. S., Kelson D. D., Reitzel D. B., Rich R. M.,
Richer H. B., 2008, ApJ, 676, 594
Kawka A., Vennes S., Schmidt G. D., Wickramasinghe D. T., Koch R., 2007,
ApJ, 654, 499
Kemic S. B., 1974, ApJ, 193, 213
Kendall T. R. et al., 2007, A&A, 466, 1059
Kilic M. et al., 2006, AJ, 131, 582
This paper has been typeset from a TEX/LATEX file prepared by the author.
C 2010 RAS, MNRAS 410, 705–716
2010 The Authors. Journal compilation 
Was this manual useful for you? yes no
Thank you for your participation!

* Your assessment is very important for improving the work of artificial intelligence, which forms the content of this project

Download PDF