to appear in the Astronomical Journal
The Solar Neighborhood XIX:
Discovery and Characterization of 33 New Nearby White Dwarf
Systems
John P. Subasavage, Todd J. Henry
Georgia State University, Atlanta, GA 30302-4106
subasavage@chara.gsu.edu, thenry@chara.gsu.edu
P. Bergeron, P. Dufour
Département de Physique, Université de Montréal, C.P. 6128, Succ. Centre-Ville,
Montréal, Québec H3C 3J7, Canada
bergeron@astro.umontreal.ca, dufourpa@astro.umontreal.ca
Nigel C. Hambly
Scottish Universities Physics Alliance (SUPA), Institute for Astronomy, University of
Edinburgh Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ, Scotland, UK
nch@roe.ac.uk
Thomas D. Beaulieu
Georgia State University, Atlanta, GA 30302-4106
beaulieu@chara.gsu.edu
ABSTRACT
We present spectra for 33 previously unclassified white dwarf systems brighter
than V = 17 primarily in the southern hemisphere. Of these new systems, 26 are
DA, 4 are DC, 2 are DZ, and 1 is DQ. We suspect three of these systems are unre-
solved double degenerates. We obtained V RI photometry for these 33 objects as
well as for 23 known white dwarf systems without trigonometric parallaxes, also
primarily in the southern hemisphere. For the 56 objects, we converted the pho-
tometry values to fluxes and fit them to a spectral energy distribution using the
http://arxiv.org/abs/0704.0894v1
– 2 –
spectroscopy to determine which model to use (i.e. pure hydrogen, pure helium,
or metal-rich helium), resulting in estimates of Teff and distance. Eight of the
new and 12 known systems are estimated to be within the NStars and Catalogue
of Nearby Stars (CNS) horizons of 25 pc, constituting a potential 18% increase
in the nearby white dwarf sample. Trigonometric parallax determinations are
underway via CTIOPI for these 20 systems.
One of the DCs is cool so that it displays absorption in the near infrared.
Using the distance determined via trigonometric parallax, we are able to constrain
the model-dependent physical parameters and find that this object is most likely a
mixed H/He atmosphere white dwarf similar to other cool white dwarfs identified
in recent years with significant absorption in the infrared due to collision-induced
absorptions by molecular hydrogen.
Subject headings: solar neighborhood — white dwarfs — stars: evolution —
stars: distances — stars: statistics
1. Introduction
The study of white dwarfs (WDs) provides insight to understanding WD formation
rates, evolution, and space density. Cool WDs, in particular, provide limits on the age of the
Galactic disk and could represent some unknown fraction of the Galactic halo dark matter.
Individually, nearby WDs are excellent candidates for astrometric planetary searches because
the astrometric signature is greater than for an identical WD system more distant. As a
population, a complete volume limited sample is necessary to provide unbiased statistics;
however, their intrinsic faintness has allowed some to escape detection.
Of the 18 WDs with trigonometric parallaxes placing them within 10 pc of the Sun (the
RECONS sample), all but one have proper motions greater than 1.′′0 yr−1 (94%). By com-
parison, of the 230 main sequence systems (as of 01 January 2007) in the RECONS sample,
50% have proper motions greater than 1.′′0 yr−1. We have begun an effort to reduce this
apparent selection bias against slower-moving WDs to complete the census of nearby WDs.
This effort includes spectroscopic, photometric, and astrometric initiatives to characterize
newly discovered as well as known WDs without trigonometric parallaxes. Utilizing the Su-
perCOSMOS Sky Survey (SSS) for plate magnitude and proper motion information coupled
with data from other recently published proper motion surveys (primarily in the southern
hemisphere), we have identified relatively bright WD candidates via reduced proper motion
diagrams.
– 3 –
In this paper, we present spectra for 33 newly discovered WD systems brighter than
V = 17.0. Once an object is spectroscopically confirmed to be a WD (in this paper for
the first time or elsewhere in the literature), we obtain CCD photometry to derive Teff and
estimate its distance using a spectral energy distribution (SED) fit and a model atmosphere
analysis. If an object’s distance estimate is within the NStars (Henry et al. 2003) and CNS
(Gliese & Jahreiß 1991) horizons of 25 pc, it is then added to CTIOPI (Cerro Tololo Inter-
American Observatory Parallax Investigation) to determine its true distance (e.g. Jao et al.
2005, Henry et al. 2006).
2. Candidate Selection
We used recent high proper motion (HPM) surveys (Pokorny et al. 2004; Subasavage et al.
2005a,b; Finch et al. 2007) in the southern hemisphere for this work because our long-term
astrometric observing program CTIOPI, is based in Chile. To select good WD candidates
for spectroscopic observations, plate magnitudes via SSS and 2MASS JHKS are extracted
for HPM objects. Each object’s (R59F − J) color and reduced proper motion (RPM) are
then plotted. RPM correlates proper motion with proximity, which is certainly not always
true; however, it is effective at separating WDs from subdwarfs and main sequence stars.
Figure 1 displays an RPM diagram for the 33 new WDs presented here. To serve as examples
for the locations of subdwarfs and main sequence stars, recent HPM discoveries from the
SuperCOSMOS-RECONS (SCR) proper motion survey are also plotted (Subasavage et al.
2005a,b). The solid line represents a somewhat arbitrary cutoff separating subdwarfs and
WDs. Targets are selected from the region below the solid line. Note there are four stars
below this line that are not represented with asterisks. Three have recently been spectro-
scopically confirmed as WDs (Subasavage et al., in preparation) and one as a subdwarf (SCR
1227−4541, denoted by “sd”) that fell just below the line at (R59F − J) = 1.4 and HR59F =
19.8 (Subasavage et al. 2005b).
Completeness limits (S/N > 10) for 2MASS are J = 15.8, H = 15.1, and KS = 14.3
for uncontaminated point sources (Skrutskie et al. 2006). The use of J provides a more
reliable RPM diagram color for objects more than a magnitude fainter than the KS limit,
which is particularly important for the WDs (with (J − KS) < 0.4) discussed here. Only
objects bright enough to have 2MASS magnitudes are included in Figure 1. Consequently,
all WD candidates are brighter than V ∼ 17, and are therefore likely to be nearby. Objects
that fall in the WD region of the RPM diagram were cross-referenced with SIMBAD and
– 4 –
McCook & Sion (1999)1 to determine those that were previously classified as WDs. The
remainder were targeted for spectroscopic confirmation.
The remaining 33 candidates comprise the “new sample” whose spectra are presented
in this work, while the “known sample” constitutes the 23 previously identified WD systems
without trigonometric parallaxes for which we have complete V RIJHKS data.
3. Data and Observations
3.1. Astrometry and Nomenclature
The traditional naming convention for WDs uses the object’s epoch 1950 equinox 1950
coordinates. Coordinates for the new sample were extracted from 2MASS along with the
Julian date of observation. These coordinates were adjusted to account for proper motion
from the epoch of 2MASS observation to epoch 2000 (hence epoch 2000 equinox 2000).
The coordinates were then transformed to equinox 1950 coordinates using the IRAF proce-
dure precess. Finally, the coordinates were again adjusted (opposite the direction of proper
motion) to obtain epoch 1950 equinox 1950 coordinates.
Proper motions were taken from various proper motion surveys in addition to unpub-
lished values obtained via the SCR proper motion survey while recovering previously known
HPM objects. Appendix A contains the proper motions used for coordinate sliding as well
as J2000 coordinates and alternate names.
3.2. Spectroscopy
Spectroscopic observations were taken on five separate observing runs in 2003 Octo-
ber and December, 2004 March and September, and 2006 May at the Cerro Tololo Inter-
American Observatory (CTIO) 1.5m telescope as part of the SMARTS Consortium. The
Ritchey-Chrétien Spectrograph and Loral 1200×800 CCD detector were used with grating
09, providing 8.6 Å resolution and wavelength coverage from 3500 to 6900 Å. Observations
consisted of two exposures (typically 20 - 30 minutes each) to permit cosmic ray rejection,
followed by a comparison HeAr lamp exposure to calibrate wavelength for each object. Bias
subtraction, dome/sky flat-fielding, and extraction of spectra were performed using standard
IRAF packages.
1The current web based catalog can be found at http://heasarc.nasa.gov/W3Browse/all/mcksion.html
http://heasarc.nasa.gov/W3Browse/all/mcksion.html
– 5 –
A slit width of 2′′ was used for the 2003 and 2004 observing runs. Some of these data have
flux calibration problems because the slit was not rotated to be aligned along the direction
of atmospheric refraction. In conjunction with telescope “jitter”, light was sometimes lost
preferentially at the red end or the blue end for these data.
A slit width of 6′′, used for the 2006 May run, eliminated most of the flux calibration
problems even though the slit was not rotated. All observations were taken at an airmass
of less than 2.0. Within our wavelength window, the maximum atmospheric differential
refraction is less than 3′′ (Filippenko 1982). A test was performed to verify that no resolution
was lost by taking spectra of a F dwarf with sharp absorption lines from slit widths of 2′′ to
10′′ in 2′′ increments. Indeed, no resolution was lost.
Spectra for the new DA WDs with Teff ≥ 10000 K are plotted in Figure 2 while spectra
for the new DA WDs with Teff < 10000 K are plotted in Figure 3. Featureless DC spectra
are plotted in Figure 4. Spectral plots as well as model fits for unusual objects are described
in § 4.2.
3.3. Photometry
Optical V RI (Johnson V , Kron-Cousins RI) for the new and known samples was ob-
tained using the CTIO 0.9 m telescope during several observing runs from 2003 through
2006 as part of the Small and Moderate Aperture Research Telescope System (SMARTS)
Consortium. The 2048×2046 Tektronix CCD camera was used with the Tek 2 V RI filter
set2. Standard stars from Graham (1982), Bessel (1990), and Landolt (1992) were observed
each night through a range of airmasses to calibrate fluxes to the Johnson-Kron-Cousins
system and to calculate extinction corrections.
Bias subtraction and dome flat-fielding (using calibration frames taken at the beginning
of each night) were performed using standard IRAF packages. When possible, an aperture
14′′ in diameter was used to determine the stellar flux, which is consistent with the aperture
used by Landolt (1992) for the standard stars. If cosmic rays fell within this aperture, they
were removed before flux extraction. In cases of crowded fields, aperture corrections were
applied and ranged from 4′′ to 12′′ in diameter using the largest aperture possible without
including contamination from neighboring sources. Uncertainties in the optical photometry
were derived by estimating the internal night-to-night variations as well as the external errors
(i.e. fits to the standard stars). A complete discussion of the error analysis can be found in
2The central wavelengths for V , R, and I are 5475, 6425, and 8075Å respectively.
– 6 –
Henry et al. (2004). We adopt a total error of ±0.03 mag in each band. The final optical
magnitudes are listed in Table 1 as well as the number of nights each object was observed.
Infrared JHKS magnitudes and errors were extracted via Aladin from 2MASS and are
also listed in Table 1. JHKS magnitude errors are, in most cases, significantly larger than for
V RI, and the errors listed give a measure of the total photometric uncertainty (i.e. include
both global and systematic components). In cases when the magnitude error is null, the star
is near the magnitude limit of 2MASS and the photometry is not reliable.
4. Analysis
4.1. Modeling of Physical Parameters
The pure hydrogen, pure helium, and mixed hydrogen and helium model atmospheres
used to model the WDs are described at length in Bergeron et al. (2001) and references
therein, while the helium-rich models appropriate for DQ and DZ stars are described in
Dufour et al. (2005, 2007), respectively. The atmospheric parameters for each star are ob-
tained by converting the optical V RI and infrared JHKS magnitudes into observed fluxes,
and by comparing the resulting SEDs with those predicted from our model atmosphere cal-
culations. The first step is accomplished by transforming the magnitudes into average stellar
fluxes fm
received at Earth using the calibration of Holberg et al. (2006) for photon count-
ing devices. The observed and model fluxes, which depend on Teff , log g, and atmospheric
composition, are related by the equation
= 4π (R/D)2 Hm
, (1)
where R/D is the ratio of the radius of the star to its distance from Earth, and Hm
the Eddington flux, properly averaged over the corresponding filter bandpass. Our fitting
technique relies on the nonlinear least-squares method of Levenberg-Marquardt (Press et al.
1992), which is based on a steepest descent method. The value of χ2 is taken as the sum over
all bandpasses of the difference between both sides of eq. (1), weighted by the corresponding
photometric uncertainties. We consider only Teff and the solid angle to be free parameters,
and the uncertainties of both parameters are obtained directly from the covariance matrix
of the fit. In this study, we simply assume a value of log g = 8.0 for each star.
As discussed in Bergeron et al. (1997, 2001), the main atmospheric constituent — hy-
drogen or helium — is determined by comparing the fits obtained with both compositions,
or by the presence of Hα in the optical spectra. For DQ and DZ stars, we rely on the
– 7 –
procedure outlined in Dufour et al. (2005, 2007), respectively: we obtain a first estimate
of the atmospheric parameters by fitting the energy distribution with an assumed value of
the metal abundances. We then fit the optical spectrum to measure the metal abundances,
and use these values to improve our atmospheric parameters from the energy distribution.
This procedure is iterated until a self-consistent photometric and spectroscopic solution is
achieved.
The derived values for Teff for each object are listed in Table 1. Also listed are the
spectral types for each object determined based on their spectral features. The DAs have
been assigned a half-integer temperature index as defined by McCook & Sion (1999), where
the temperature index equals 50,400/Teff. As an external check, we compare in Figure 5
the photometric effective temperatures for the DA stars in Table 1 with those obtained by
fitting the observed Balmer line profiles (Figs. 2 and 3) using the spectroscopic technique
developed by Bergeron et al. (1992b), and recently improved by Liebert et al. (2003). Our
grid of pure hydrogen, NLTE, and convective model atmospheres is also described in Liebert
et al. The uncertainties of the spectroscopic technique are typically of 0.038 dex in log g
and 1.2% in Teff according to that study. We adopt a slightly larger uncertainty of 1.5%
in Teff (Spec) because of the problematic flux calibrations of the pre−2006 data (see § 3.2).
The agreement shown in Figure 5 is excellent, except perhaps at high temperatures where
the photometric determinations become more uncertain. It is possible that the significantly
elevated point in Figure 5, WD 0310−624 (labeled), is an unresolved double degenerate (see
§ 4.2). We refrain here from using the log g determinations in our analysis because these
are available only for the DA stars in our sample, and also because the spectra are not flux
calibrated accurately enough for that purpose.
Once the effective temperature and the atmospheric composition are determined, we
calculate the absolute visual magnitude of each star by combining the new calibration of
Holberg et al. (2006) with evolutionary models similar to those described in Fontaine et al.
(2001) but with C/O cores, q(He) ≡ logMHe/M⋆ = 10
−2 and q(H) = 10−4 (representative
of hydrogen-atmosphere WDs), and q(He) = 10−2 and q(H) = 10−10 (representative of
helium-atmosphere WDs)3. By combining the absolute visual magnitude with the Johnson
V magnitude, we derive a first estimate of the distance of each star (reported in Table 1).
Errors on the distance estimates incorporate the errors of the photometry values as well as
an error of 0.25 dex in log g, which is the measured dispersion of the observed distribution
using spectroscopic determinations (see Figure 9 of Bergeron et al. 1992b).
Of the 33 new systems presented here, 5 have distance estimates within 25 pc. Four
3see http://www.astro.umontreal.ca/˜bergeron/CoolingModels/
http://www.astro.umontreal.ca/~bergeron/CoolingModels/
– 8 –
more systems require additional attention because distance estimates are derived via other
means. Three of these are likely within 25 pc. All four are further discussed in the next
section. In total, 20 WD systems (8 new and 12 known) are estimated (or determined) to be
within 25 pc and one additional common proper motion binary system possibly lies within
25 pc.
4.2. Comments on Individual Systems
Here we address unusual and interesting objects.
WD 0121−429 is a DA WD that exhibits Zeeman splitting of Hα and Hβ, thereby
making its formal classification DAH. The SED fit to the photometry is superb, yielding
a Teff of 6,369 ± 137 K. When we compare the strength of the absorption line trio with
that predicted using the Teff from the SED fit, the depth of the absorption appears too
shallow. Using the magnetic line fitting procedure outlined in Bergeron et al. (1992a), we
must include a 50% dilution factor to match the observed central line of Hα. In light of
this, we utilized the trigonometric parallax distance determined via CTIOPI of 17.7 ± 0.7
pc (Subasavage et al., in preparation) to further constrain this system. The resulting SED
fit, with distance (hence luminosity) as a constraint rather than a variable, implies a mass of
0.43 ± 0.03 M⊙. Given the age of our Galaxy, the lowest mass WD that could have formed
is ∼0.47 M⊙ (Iben & Renzini 1984). It is extremely unlikely that this WD formed through
single star evolution. The most likely scenario is that this is a double degenerate binary
with a magnetic DA component and a featureless DC component (necessary to dilute the
absorption at Hα), similar to G62-46 (Bergeron et al. 1993) and LHS 2273 (see Figure 33
of Bergeron et al. 1997). If this interpretation is correct, any number of component masses
and luminosities can reproduce the SED fit.
The spectrum and corresponding magnetic fit to the Hα lines (including the dilution)
is shown in Figure 6. The viewing angle, i = 65◦, is defined as the angle between the dipole
axis and the line of sight (i = 0 corresponds to a pole-on view). The best fit produces a
dipole field strength, Bd = 9.5 MG, and a dipole offset, az = 0.06 (in units of stellar radius).
The positive value of az implies that the offset is toward the observer. Only Bd is moderately
constrained, both i and az can vary significantly yet still produce a reasonable fit to the data
(Bergeron et al. 1992a).
WD 0310−624 is a DAWD that is one of the hottest in the new sample. Because of it’s
elevation significantly above the equal temperature line (solid) in Figure 5, it is possible that
it is an unresolved double degenerate with very different component effective temperatures.
– 9 –
In fact, this method has been used to identify unresolved double degenerate candidates (i.e.
Bergeron et al. 2001).
WD 0511−415 is a DA WD (spectrum is plotted in Figure 2) whose spectral fit
produces a Teff = 10,813 ± 219 K and a log g = 8.21 ± 0.10 using the spectral fitting
procedure of Liebert et al. (2003). This object lies near the red edge of the ZZ Ceti instability
strip as defined by Gianninas et al. (2006). If variable, this object would help to constrain
the cool edge of the instability strip in Teff , log g parameter space. Follow-up high speed
photometry is necessary to confirm variability.
WD 0622−329 is a DAB WD displaying the Balmer lines as well as weaker He I at
4472 and 5876 Å. The spectrum, shown in Figure 7, is reproduced best with a model having
Teff ∼43,700 K. However, the predicted He II absorption line at 4686 Å for a WD of this
Teff is not present in the spectrum. In contrast, the SED fit to the photometry implies a
Teff of ∼10,500 K (using either pure H or pure He models). Because the Teff values are
vastly discrepant, we explored the possibility that this spectrum is not characterized by a
single temperature. We modeled the spectrum assuming the object was an unresolved double
degenerate. The best fit implies one component is a DB with Teff = 14,170 ± 1,228 K and
the other component is a DA with Teff = 9,640 ± 303 K, similar to the unresolved DA +
DB degenerate binary PG 1115+166 analyzed by Bergeron & Liebert (2002). One can see
from Figure 7 that the spectrum is well modeled under this assumption. We conclude this
object is likely a distant (well beyond 25 pc) unresolved double degenerate.
WD 0840−136 is a DZ WD whose spectrum shows both Ca II (H & K) and Ca I (4226
Å) lines as shown in Figure 8. Fits to the photometric data for different atmospheric com-
positions indicate temperatures of about 4800-5000 K. However, fits to the optical spectrum
using the models of Dufour et al. (2007) cannot reproduce simultaneously all three calium
lines. This problem is similar to that encountered by Dufour et al. (2007) where the atmo-
spheric parameters for the coolest DZ WDs were considered uncertain because of possible
high atmospheric pressure effects. We utilize a photometric relation relevant for WDs of any
atmospheric composition, which links MV to (V −I) (Salim et al. 2004) to obtain a distance
estimate of 19.3 ± 3.9 pc.
WD 1054−226 was observed spectroscopically as part of the Edinburgh-Cape (EC)
blue object survey and assigned a spectral type of sdB+ (Kilkenny et al. 1997). As is evident
in Figure 3, the spectrum of this object is the noisiest of all the spectra presented here and
perhaps a bit ambiguous. As an additional check, this object was recently observed using
the ESO 3.6 m telescope and has been confirmed to be a cool DA WD (Bergeron, private
communication).
– 10 –
WD 1105−340 is a DA WD (spectrum is plotted in Figure 2) with a common proper
motion companion with separation of 30.′′6 at position angle 107.1◦. The companion’s spectral
type is M4Ve with VJ = 15.04, RKC = 13.68, IKC = 11.96, J = 10.26, H = 9.70, and KS
= 9.41. In addition to the SED derived distance estimate for the WD, we utilize the main
sequence distance relations of Henry et al. (2004) to estimate a distance to the red dwarf
companion. We obtain a distance estimate of 19.1 ± 3.0 pc for the companion leaving
open the possibility that this system may lie just within 25 pc. A trigonometric parallax
determination is currently underway for confirmation.
WD 1149−272 is the only DQ WD discovered in the new sample. This object was
observed spectroscopically as part of the Edinburgh-Cape (EC) blue object survey for which
no features deeper than 5% were detected and was labeled a possible DC (Kilkenny et al.
1997). It is identified as having weak C2 swan band absorption at 4737 and 5165 Å and is
otherwise featureless. The DQ model reproduces the spectrum reliably and is overplotted in
Figure 9. This object is characterized as having Teff = 6188 ± 194 K and a log (C/He) =
−7.20 ± 0.16.
WD 2008−600 is a DC WD (spectrum is plotted in Figure 4) that is flux deficient in
the near infrared, as indicated by the 2MASS magnitudes. The SED fit to the photometry
is a poor match to either the pure hydrogen or the pure helium models. A pure hydrogen
model provides a slightly better match than a pure helium model, and yields a Teff of ∼3100
K, thereby placing it in the relatively small sample of ultracool WDs. In order to discern the
true nature of this object, we have constrained the model using the distance obtained from
the CTIOPI trigonometric parallax of 17.1 ± 0.4 pc (Subasavage et al., in preparation). This
object is then best modeled as having mostly helium with trace amounts of hydrogen (log
(He/H) = 2.61) in its atmosphere and has a Teff = 5078 ± 221 K (see Figure 10). A mixed
hydrogen and helium composition is required to produce sufficient absorption in the infrared
as a result of the collision-induced absorption by molecular hydrogen due to collisions with
helium. Such mixed atmospheric compositions have also been invoked to explain the infrared
flux deficiency in LHS 1126 (Bergeron et al. 1994) as well as SDSS 1337+00 and LHS 3250
(Bergeron & Leggett 2002). While WD 2008−600 is likely not an ultracool WD, it is one of
the brightest and nearest cool WDs known. Because the 2MASS magnitudes are not very
reliable, we intend to obtain additional near-infrared photometry to better constrain the fit.
WD 2138−332 is a DZ WD for which a calcium rich model reproduces the spectrum
reliably. The spectrum and the overplotted fit are shown in the bottom panel of Figure
8. Clearly evident in the spectrum are the strong Ca II absorption at 3933 and 3968 Å. A
weaker Ca I line is seen at 4226Å. Also seen are Mg I absorption lines at 3829, 3832, and
3838 Å (blended) as well as Mg I at 5167, 5173, and 5184 Å (also blended). Several weak Fe I
– 11 –
lines from 4000Å to 4500Å and again from 5200Å to 5500Å are also present. The divergence
of the spectrum from the fit toward the red end is likely due to an imperfect flux calibration
of the spectrum. This object is characterized as having Teff = 7188 ± 291 K and a log
(Ca/He) = −8.64 ± 0.16. The metallicity ratios are, at first, assumed to be solar (as defined
by Grevesse & Sauval 1998) and, in this case, the quality of the fit was sufficient without
deviation. The corresponding log (Mg/He) = −7.42 ± 0.16 and log (Fe/He) = −7.50 ± 0.16
for this object.
WD 2157−574 is a DAWD (spectrum is plotted in Figure 3) unique to the new sample
in that it displays weak Ca II absorption at 3933 and 3968 Å (H and K) thereby making its
formal classification a DAZ. Possible scenarios that enrich the atmospheres of DAZs include
accretion via (1) debris disks, (2) ISM, and (3) cometary impacts (see Kilic et al. 2006 and
references therein). The 2MASS KS magnitude is near the faint limit and is unreliable, but
even considering the J and H magnitudes, there appears to be no appreciable near-infrared
excess. While this may tentatively rule out the possibility of a debris disk, this object would
be an excellent candidate for far-infrared spaced-based studies to ascertain the origin of the
enrichment.
5. Discussion
WDs represent the end state for stars less massive than ∼8 M⊙ and are therefore rel-
atively numerous. Because of their intrinsic faintness, only the nearby WD population can
be easily characterized and provides the benchmark upon which WD stellar astrophysics is
based. It is clear from this work and others (e.g. Holberg et al. 2002; Kawka & Vennes 2006)
that the WD sample is complete, at best, to only 13 pc. Spectroscopic confirmation of new
WDs as well as trigonometric parallax determinations for both new and known WDs will
lead to a more complete sample and will push the boundary of completeness outward. We
estimate that 8 new WDs and an additional 12 known WDs without trigonometric parallaxes
are nearer than 25 pc, including one within 10 pc (WD 0141−675). Parallax measurements
via CTIOPI are underway for these 20 objects to confirm proximity. This total of 20 WDs
within 25 pc constitutes an 18% increase to the 109 WDs with trigonometric parallaxes ≥
40 mas.
Evaluating the proper motions of the new and known samples within 25 pc indicates
that almost double the number of systems have been found with µ < 1.′′0 yr−1 than with
µ ≥ 1.′′0 yr−1 (13 vs 7, see Table 2). The only WD estimated to be within 10 pc has µ >
1.′′0 yr−1, although WD 1202−232 is estimated to be 10.2 ± 1.7 pc and it’s proper motion is
small (µ = 0.′′227 yr−1).
– 12 –
Because this effort focuses mainly on the southern hemisphere, it is likely that there
is a significant fraction of nearby WDs in the northern hemisphere that have also gone
undetected. With the recent release of the LSPM-North Catalog (Lépine & Shara 2005),
these objects are identifiable by employing the same techniques used in this work. The
challenge is the need for a large scale parallax survey focusing on WDs to confirm proximity.
Since the HIPPARCOS mission, only six WD trigonometric parallaxes have been published
(Hambly et al. 1999; Smart et al. 2003), and of those, only two are within 25 pc. The USNO
parallax program is in the process of publishing trigonometric parallaxes for ∼130 WDs,
mostly in the northern hemisphere, although proximity was not a primary motivation for
target selection (Dahn, private communication).
In addition to further completing the nearby WD census, the wealth of observational
data available from this effort provides reliable constraints on their physical parameters
(i.e. Teff , log g, mass, and radius). Unusual objects are then revealed, such as those dis-
cussed in § 4.2. In particular, trigonometric parallaxes help identify WDs that are overlu-
minous, as is the case for WD 0121−429. This object, and others similar to it, are excellent
candidates to provide insight into binary evolution. If they can be resolved using high res-
olution astrometric techniques (i.e. speckle, adaptive optics, or interferometry via Hubble
Space Telescope’s Fine Guidance Sensors), they may provide astrometric masses, which are
fundamental calibrators for stellar structure theory and for the reliability of the theoretical
WD mass-radius and initial-to-final-mass relationships. To date, only four WD astrometric
masses are known to better than ∼ 5% (Provencal et al. 1998).
One avenue that is completely unexplored to date is a careful high resolution search
for planets around WDs. Theory dictates that the Sun will become a WD, and when it
does, the outer planets will remain in orbit (not without transformations of their own, of
course). In this scenario, the Sun will have lost more than half of its mass, thereby amplifying
the signature induced by the planets. Presumably, this has already occurred in the Milky
Way and systems such as these merely await detection. Because of the faintness and spectral
signatures of WDs (i.e. few, if any, broad absorption lines), current radial velocity techniques
are inadequate for planet detection, leaving astrometric techniques as the only viable option.
For a given system, the astrometric signature is inversely related to distance (i.e. the nearer
the system, the larger the astrometric signature). This effort aims to provide a complete
census of nearby WDs that can be probed for these astrometric signatures using future
astrometric efforts.
– 13 –
6. Acknowledgments
The RECONS team at Georgia State University wishes to thank the NSF (grant AST
05-07711), NASA’s Space Interferometry Mission, and GSU for their continued support of
our study of nearby stars. We also thank the continuing support of the members of the
SMARTS consortium, who enable the operations of the small telescopes at CTIO where all
of the data in this work were collected. J. P. S. is indebted to Wei-Chun Jao for the use of his
photometry reduction pipeline. P. B. is a Cottrell Scholar of Research Corporation and would
like to thank the NSERC Canada for its support. N. C. H. would like to thank colleagues in
the Wide Field Astronomy Unit at Edinburgh for their efforts contributing to the existence
of the SSS; particular thanks go to Mike Read, Sue Tritton, and Harvey MacGillivray. This
work has made use of the SIMBAD, VizieR, and Aladin databases, operated at the CDS in
Strasbourg, France. We have also used data products from the Two Micron All Sky Survey,
which is a joint project of the University of Massachusetts and the Infrared Processing and
Analysis Center, funded by NASA and NSF.
A. Appendix
In order to ensure correct cross-referencing of names for the new and known WD systems
presented here, Table 3 lists additional names found in the literature. Objects for which there
is an NLTT designation will also have the corresponding L or LP designations found in the
NLTT catalog. This is necessary because the NLTT designations were not published in the
original catalog, but rather are the record numbers in the electronic version of the catalog
and have been adopted out of necessity.
REFERENCES
Bergeron, P., Ruiz, M.-T., & Leggett, S. K. 1992, ApJ, 400, 315
Bergeron, P., Saffer, R. A., & Liebert, J. 1992, ApJ, 394, 228
Bergeron, P., Ruiz, M.-T., & Leggett, S. K. 1993, ApJ, 407, 733
Bergeron, P., Ruiz, M.-T., Leggett, S. K., Saumon, D., & Wesemael, F. 1994, ApJ, 423, 456
Bergeron, P., Ruiz, M. T., & Leggett, S. K. 1997, ApJS, 108, 339
Bergeron, P., Leggett, S. K., & Ruiz, M. T. 2001, ApJS, 133, 413
– 14 –
Bergeron, P., & Leggett, S. K. 2002, ApJ, 580, 1070
Bergeron, P., & Liebert, J. 2002, ApJ, 566, 1091
Bessel, M. S. 1990, A&AS, 83, 357
Carpenter, J. M. 2001, AJ, 121, 2851
Dufour, P., Bergeron, P., & Fontaine, G. 2005, ApJ, 627, 404
Dufour, P., Bergeron, P., Liebert, J., Harris, H. C., Knapp, G. R., Anderson, S. F., Hall,
P. B., Strauss, M. A., Collinge, M. J., & Edwards, M. C. 2007, submitted to ApJ
Filippenko, A. V. 1982, PASP, 94, 715
Finch, C. T., Henry, T. J., Subasavage, J. P., Jao, W.-C., Hambly, N. C. 2007, AJ, submitted
Fontaine, G., Brassard, P., & Bergeron, P. 2001, PASP, 113, 409
Gianninas, A., Bergeron, P., & Fontaine, G. 2006, AJ, 132, 831
Gliese, W., & Jahreiß, H. 1991, On: The Astronomical Data Center CD-ROM: Selected As-
tronomical Catalogs, Vol. I; L.E. Brotzmann, S.E. Gesser (eds.), NASA/Astronomical
Data Center, Goddard Space Flight Center, Greenbelt, MD
Graham, J. A. 1982, PASP, 94, 244
Grevesse, N., & Sauval, A. J. 1998, Space Science Reviews, 85, 161
Hambly, N. C., Smartt, S. J., Hodgkin, S. T., Jameson, R. F., Kemp, S. N., Rolleston,
W. R. J., & Steele, I. A. 1999, MNRAS, 309, L33
Henry, T. J., Walkowicz, L. M., Barto, T. C., & Golimowski, D. A. 2002, AJ, 123, 2002
Henry, T. J., Backman, D. E., Blackwell, J., Okimura, T., & Jue, S. 2003, The Future of
Small Telescopes In The New Millennium. Volume III - Science in the Shadow of
Giants, 111
Henry, T. J., Subasavage, J. P., Brown, M. A., Beaulieu, T. D., Jao, W., & Hambly, N. C.
2004, AJ, 128, 2460
Henry, T. J., Jao, W.-C., Subasavage, J. P., Beaulieu, T. D., Ianna, P. A., Costa, E., &
Méndez, R. A. 2006, AJ, 132, 2360
Holberg, J. B., Oswalt, T. D., & Sion, E. M. 2002, ApJ, 571, 512
– 15 –
Holberg, J. B., & Bergeron, P. 2006, AJ, 132, 1223
Iben, I., & Renzini, A. 1984, Phys. Rep., 105, 329
Jao, W.-C., Henry, T. J., Subasavage, J. P., Brown, M. A., Ianna, P. A., Bartlett, J. L.,
Costa, E., & Méndez, R. A. 2005, AJ, 129, 1954
Kawka, A., & Vennes, S. 2006, ApJ, 643, 402
Kilic, M., von Hippel, T., Leggett, S. K., & Winget, D. E. 2006, ApJ, 646, 474
Kilkenny, D., O’Donoghue, D., Koen, C., Stobie, R. S., & Chen, A. 1997, MNRAS, 287, 867
Kleinman, S. J., et al. 2004, ApJ, 607, 426
Landolt, A. U. 1992, AJ, 104, 340
Lépine, S., Shara, M. M., & Rich, R. M. 2003, AJ, 126, 921
Lépine, S., & Shara, M. M. 2005, AJ, 129, 1483
Lépine, S., Rich, R. M., & Shara, M. M. 2005, ApJ, 633, L121
Liebert, J., Bergeron, P., & Holberg, J. B. 2003, AJ, 125, 348
Liebert, J., Bergeron, P., & Holberg, J. B. 2005, ApJS, 156, 47
Luyten, W. J. 1949, ApJ, 109, 528
Luyten, W. J. 1979, LHS Catalogue (2nd ed.; Minneapolis: Univ. of Minnesota Press)
Luyten, W. J. 1979, New Luyten Catalogue of Stars with Proper Motions Larger than Two
Tenths of an Arcsecond (Minneapolis: Univ. of Minnesota Press)
McCook, G. P., & Sion, E. M. 1999, ApJS, 121, 1
Oppenheimer, B. R., Hambly, N. C., Digby, A. P., Hodgkin, S. T., & Saumon, D. 2001,
Science, 292, 698
Press, W. H., Teukolsky, S. A., Vetterling, W. T., & Flannery, B. P. 1992, Numerical Recipes
in FORTRAN, 2nd edition (Cambridge: Cambridge University Press), 644
Provencal, J. L., Shipman, H. L., Hog, E., & Thejll, P. 1998, ApJ, 494, 759
Pokorny, R. S., Jones, H. R. A., Hambly, N. C., & Pinfield, D. J. 2004, A&A, 421, 763
– 16 –
Salim, S., Rich, R. M., Hansen, B. M., Koopmans, L. V. E., Oppenheimer, B. R., & Bland-
ford, R. D. 2004, ApJ, 601, 1075
Scholz, R.-D., Szokoly, G. P., Andersen, M., Ibata, R., & Irwin, M. J. 2002, ApJ, 565, 539
Skrutskie, M. F., et al. 2006, AJ, 131, 1163
Smart, R. L., et al. 2003, A&A, 404, 317
Subasavage, J. P., Henry, T. J., Hambly, N. C., Brown, M. A., & Jao, W. 2005, AJ, 129, 413
Subasavage, J. P., Henry, T. J., Hambly, N. C., Brown, M. A., Jao, W.-C., & Finch, C. T.
2005, AJ, 130, 1658
This preprint was prepared with the AAS LATEX macros v5.2.
Table 1. Optical and Infrared Photometry, and Derived Parameters for New and Known White Dwarfs.
WD VJ RC IC # J σJ H σH KS σK
Teff Comp Dist SpT Notes
Name Obs (K) (pc)
New Spectroscopically Confirmed White Dwarfs
0034−602............. 14.08 14.19 14.20 3 14.37 0.04 14.55 0.06 14.52 0.09 14655±1413 H 35.8±5.7 DA3.5
0121−429............. 14.83 14.52 14.19 4 13.85 0.02 13.63 0.04 13.53 0.04 6369± 137 H · · · ± · · · DAH a
0216−398............. 15.75 15.55 15.29 3 15.09 0.04 14.83 0.06 14.89 0.14 7364± 241 H 29.9±4.7 DA7.0
0253−755............. 16.70 16.39 16.08 2 15.77 0.07 15.76 0.15 15.34 null 6235± 253 He 34.7±5.5 DC
0310−624............. 15.92 15.99 16.03 2 16.13 0.10 16.31 0.27 16.50 null 13906±1876 H · · · ± · · · DA3.5 b
0344+014............. 16.52 16.00 15.54 2 15.00 0.04 14.87 0.09 14.70 0.12 5084± 91 He 19.9±3.1 DC
0404−510............. 15.81 15.76 15.70 2 15.74 0.06 15.55 0.13 15.59 null 10052± 461 H 53.5±8.5 DA5.0
0501−555............. 16.35 16.17 15.98 2 15.91 0.08 15.72 0.15 15.82 0.26 7851± 452 He 44.8±6.9 DC
0511−415............. 16.00 15.99 15.93 2 15.96 0.08 15.97 0.15 15.20 null 10393± 560 H 61.8±10.8 DA5.0
0525−311............. 15.94 16.03 16.03 2 16.20 0.12 16.21 0.25 14.98 null 12941±1505 H 76.3±13.6 DA4.0
0607−530............. 15.99 15.92 15.78 3 15.82 0.07 15.66 0.14 15.56 0.21 9395± 426 H 51.7±9.0 DA5.5
0622−329............. 15.47 15.41 15.36 2 15.44 0.06 15.35 0.11 15.53 0.25 · · · ± · · · · · · · · · ± · · · DAB c
0821−669............. 15.34 14.82 14.32 3 13.79 0.03 13.57 0.03 13.34 0.04 5160± 95 H 11.5±1.9 DA10.0
0840−136............. 15.72 15.36 15.02 3 14.62 0.03 14.42 0.05 14.54 0.09 · · · ± · · · · · · · · · ± · · · DZ d
1016−308............. 14.67 14.75 14.81 2 15.05 0.04 15.12 0.08 15.41 0.21 16167±1598 H 50.6±9.2 DA3.0
1054−226............. 16.02 15.82 15.62 2 15.52 0.05 15.40 0.11 15.94 0.26 8266± 324 H 41.0±7.0 DA6.0 e
1105−340............. 13.66 13.72 13.79 2 13.95 0.03 13.98 0.04 14.05 0.07 13926± 988 H 28.2±4.8 DA3.5 f
1149−272............. 15.87 15.59 15.37 4 15.17 0.05 14.92 0.06 14.77 0.11 6188± 194 He (+C) 24.0±3.8 DQ
1243−123............. 15.57 15.61 15.64 2 15.74 0.07 15.73 0.11 16.13 null 12608±1267 H 62.6±10.7 DA4.0
1316−215............. 16.67 16.33 15.99 2 15.56 0.05 15.33 0.08 15.09 0.14 6083± 201 H 31.6±5.3 DA8.5
1436−781............. 16.11 15.82 15.49 2 15.04 0.04 14.88 0.08 14.76 0.14 6246± 200 H 26.0±4.3 DA8.0
1452−310............. 15.85 15.77 15.63 2 15.58 0.06 15.54 0.09 15.50 0.22 9206± 375 H 46.8±8.1 DA5.5
1647−327............. 16.21 15.85 15.49 3 15.15 0.05 14.82 0.08 14.76 0.11 6092± 193 H 25.5±4.2 DA8.5
1742−722............. 15.53 15.62 15.70 2 15.85 0.08 15.99 0.18 15.65 null 15102±2451 H 71.7±12.9 DA3.5
1946−273............. 14.19 14.31 14.47 2 14.72 0.04 14.77 0.09 14.90 0.13 21788±3304 H 52.0±9.9 DA2.5
2008−600............. 15.84 15.40 14.99 4 14.93 0.05 15.23 0.11 15.41 null 5078± 221 He · · · ± · · · DC g
2008−799............. 16.35 15.96 15.57 3 15.11 0.04 15.03 0.08 14.64 0.09 5807± 161 H 24.5±4.1 DA8.5
2035−369............. 14.94 14.85 14.72 2 14.75 0.04 14.72 0.06 14.84 0.09 9640± 298 H 33.1±5.7 DA5.0
2103−397............. 15.31 15.15 14.91 2 14.79 0.03 14.63 0.04 14.64 0.08 7986± 210 H 28.2±4.8 DA6.5
2138−332............. 14.47 14.30 14.16 3 14.17 0.03 14.08 0.04 13.95 0.06 7188± 291 He (+Ca) 17.3±2.7 DZ
2157−574............. 15.96 15.73 15.49 3 15.18 0.04 15.05 0.07 15.28 0.17 7220± 246 H 32.0±5.4 DAZ
2218−416............. 15.36 15.35 15.24 2 15.38 0.04 15.14 0.09 15.39 0.15 10357± 414 H 45.6±8.0 DA5.0
2231−387............. 16.02 15.88 15.62 2 15.57 0.06 15.51 0.11 15.11 0.15 8155± 336 H 40.6±6.9 DA6.0
Known White Dwarfs without a Trigonometric Parallax Estimated to be Within 25 pc
0141−675 ............ 13.82 13.52 13.23 3 12.87 0.02 12.66 0.03 12.58 0.03 6484± 128 H 9.7±1.6 DA8.0
0806−661 ............ 13.73 13.66 13.61 3 13.70 0.02 13.74 0.03 13.78 0.04 10753± 406 He 21.1±3.5 DQ
1009−184 ............ 15.44 15.18 14.91 3 14.68 0.04 14.52 0.05 14.31 0.07 6449± 194 He 20.9±3.2 DZ h
1036−204 ............ 16.24 15.54 15.34 3 14.63 0.03 14.35 0.04 14.03 0.07 4948± 70 He 16.2±2.5 DQ i
1202−232 ............ 12.80 12.66 12.52 3 12.40 0.02 12.30 0.03 12.34 0.03 8623± 168 H 10.2±1.7 DA6.0
1315−781 ............ 16.16 15.73 15.35 2 14.89 0.04 14.67 0.08 14.58 0.12 5720± 162 H 21.6±3.6 DC j
1339−340 ............ 16.43 16.00 15.56 2 15.00 0.04 14.75 0.06 14.65 0.10 5361± 138 H 21.2±3.5 DA9.5
1756+143 ............ 16.30 16.12 15.69 1 14.93 0.04 14.66 0.06 14.66 0.08 5466± 151 H 22.4±3.4 DA9.0 k
Table 1—Continued
WD VJ RC IC # J σJ H σH KS σK
Teff Comp Dist SpT Notes
Name Obs (K) (pc)
1814+134 ............ 15.85 15.34 14.86 2 14.38 0.04 14.10 0.06 14.07 0.06 5313± 115 H 15.6±2.5 DA9.5
2040−392 ............ 13.74 13.77 13.68 2 13.77 0.02 13.82 0.03 13.81 0.05 10811± 325 H 23.1±4.0 DA4.5
2211−392 ............ 15.91 15.61 15.24 2 14.89 0.03 14.64 0.05 14.56 0.08 6243± 167 H 23.5±4.0 DA8.0
2226−754A........... 16.57 15.93 15.33 2 14.66 0.04 14.66 0.06 14.44 0.08 4230± 104 H 12.8±2.0 DC l
2226−754B........... 16.88 16.17 15.51 2 14.86 0.04 14.82 0.06 14.72 0.12 4177± 112 H 14.0±2.2 DC l
Known White Dwarfs without a Trigonometric Parallax Estimated to be Beyond 25 pc
0024−556............. 15.17 15.15 15.07 2 15.01 0.04 15.23 0.10 15.09 0.14 10007± 378 H 39.8±6.8 DA5.0
0150+256............. 15.70 15.52 15.33 2 15.07 0.04 15.07 0.09 15.15 0.14 7880± 280 H 33.0±5.6 DA6.5
0255−705 ............ 14.08 14.03 14.00 2 14.04 0.03 14.12 0.04 13.99 0.06 10541± 326 H 25.8±4.5 DA5.0
0442−304............. 16.03 15.93 15.86 2 15.94 0.09 15.81 null 15.21 null 9949± 782 He 55.1±9.1 DQ
0928−713 ............ 15.11 14.97 14.83 3 14.77 0.03 14.69 0.06 14.68 0.09 8836± 255 H 30.7±5.3 DA5.5
1143−013............. 16.39 16.08 15.79 1 15.54 0.06 15.38 0.08 15.18 0.16 6824± 250 H 34.4±5.8 DA7.5
1237−230 ............ 16.53 16.13 15.74 2 15.35 0.05 15.08 0.08 14.94 0.11 5841± 173 H 26.9±4.5 DA8.5
1314−153............. 14.82 14.89 14.97 2 15.17 0.05 15.26 0.09 15.32 0.21 15604±2225 H 52.7±9.5 DA3.0
1418−088 ............ 15.39 15.21 15.01 2 14.76 0.04 14.73 0.06 14.76 0.10 7872± 243 H 28.5±4.8 DA6.5
1447−190............. 15.80 15.59 15.32 2 15.06 0.04 14.87 0.07 14.78 0.11 7153± 235 H 29.1±4.9 DA7.0
1607−250............. 15.19 15.12 15.09 2 15.08 0.08 15.08 0.08 15.22 0.15 10241± 457 H 41.2±7.2 DA5.0
aDistance via SED fit (not listed) is underestimated because object is likely an unresolved double degenerate with one magnetic component (see § 4.2). Instead, we
adopt the trigonometric parallax distance of 17.7 ± 0.7 pc derived via CTIOPI.
bDistance via SED fit (not listed) is underestimated because object is likely a distant (well beyond 25 pc) unresolved double degenerate (see § 4.2).
cDistance via SED fit (not listed) is underestimated because object is likely a distant (well beyond 25 pc) unresolved double degenerate with components of type DA
and DB (see § 4.2). Temperatures derived from the spectroscopic fit yield 9,640 ± 303 K and 14,170 ± 1,228 K for the DA and DB respectively.
dObject is likely cooler than Teff ∼5000 K and the theoretical models do not provide an accurate treatment at these temperatures (see § 4.2). Instead, we use the
linear photometric distance relation of Salim et al. (2004) and obtain a distance estimate of 19.3 ± 3.9 pc.
eThis object was observed as part of the Edinburgh-Cape survey and was classified as a sdB+ (Kilkenny et al. 1997).
fDistance of 19.1 ± 3.0 pc is estimated using V RIJHKS for the common proper motion companion M dwarf and the relations of Henry et al. (2004). System is
possibly within 25 pc. (see § 4.2).
gDistance estimate is undetermined. Instead, we adopt the distance measured via trogonometric parallax of 17.1 ± 0.4 pc (see § 4.2).
hNot listed in McCook & Sion (1999) but identified as a DC/DQ WD by Henry et al. (2002). We obtained blue spectra that show Ca II H & K absorption and classify
this object as a DZ.
iThe SED fit to the photometry is marginal. This object displays deep swan band absorption that significantly affects its measured magnitudes.
jNot listed in McCook & Sion (1999) but identified as a WD by Luyten (1949). Spectral type is derived from our spectra.
kAs of mid-2004, object has moved onto a background source. Photometry is probably contaminated, which is consistent with the poor SED fit for this object.
lSpectral type was determined using spectra published by Scholz et al. (2002).
– 19 –
Table 2. Distance Estimate Statistics for New and Known White Dwarfs.
Proper motion d ≤ 10 pc 10 pc < d ≤ 25 pc d > 25 pc
µ ≥ 1.′′0 yr−1......................... 1 6 1
1.′′0 yr−1 > µ ≥ 0.′′8 yr−1...... 0 0 0
0.′′8 yr−1 > µ ≥ 0.′′6 yr−1...... 0 2 2
0.′′6 yr−1 > µ ≥ 0.′′4 yr−1...... 0 6 11
0.′′4 yr−1 > µ ≥ 0.′′18 yr−1.... 0 5 22
Total.................................... 1 19 36
– 20 –
Table 3. Astrometry and Alternate Designations for New and Known White Dwarfs.
WD Name RA Dec PM PA Ref Alternate Names
(J2000.0) (J2000.0) (arcsec yr−1) (deg)
New Spectroscopically Confirmed White Dwarfs
0034−602......... 00 36 22.31 −59 55 27.5 0.280 069.0 L NLTT 1993 = LP 122-4 = · · ·
0121−429......... 01 24 03.98 −42 40 38.5 0.538 155.2 L LHS 1243 = NLTT 4684 = LP 991-16
0216−398......... 02 18 31.51 −39 36 33.2 0.500 078.6 L LHS 1385 = NLTT 7640 = LP 992-99
0253−755......... 02 52 45.64 −75 22 44.5 0.496 063.5 S SCR 0252-7522 = · · · = · · ·
0310−624......... 03 11 21.34 −62 15 15.7 0.416 083.3 S SCR 0311-6215 = · · · = · · ·
0344+014......... 03 47 06.82 +01 38 47.5 0.473 150.4 S LHS 5084 = NLTT 11839 = LP 593-56
0404−510......... 04 05 32.86 −50 55 57.8 0.320 090.7 P LEHPM 1-3634 = · · · = · · ·
0501−555......... 05 02 43.43 −55 26 35.2 0.280 191.9 P LEHPM 1-3865 = · · · = · · ·
0511−415......... 05 13 27.80 −41 27 51.7 0.292 004.4 P LEHPM 2-1180 = · · · = · · ·
0525−311......... 05 27 24.33 −31 06 55.7 0.379 200.7 P NLTT 15117 = LP 892-45 = LEHPM 2-521
0607−530......... 06 08 43.81 −53 01 34.1 0.246 327.6 P LEHPM 2-2008 = · · · = · · ·
0622−329......... 06 24 25.78 −32 57 27.4 0.187 177.7 P LEHPM 2-5035 = · · · = · · ·
0821−669......... 08 21 26.70 −67 03 20.1 0.758 327.6 S SCR 0821-6703 = · · · = · · ·
0840−136......... 08 42 48.45 −13 47 13.1 0.272 263.0 S NLTT 20107 = LP 726-1 = · · ·
1016−308......... 10 18 39.84 −31 08 02.0 0.212 304.0 L NLTT 23992 = LP 904-3 = LEHPM 2-5779
1054−226......... 10 56 38.64 −22 52 55.9 0.277 349.7 P NLTT 25792 = LP 849-31 = LEHPM 2-1372
1105−340......... 11 07 47.89 −34 20 51.4 0.287 168.0 S SCR 1107-3420A = · · · = · · ·
1149−272......... 11 51 36.10 −27 32 21.0 0.199 278.3 P LEHPM 2-4051 = · · · = · · ·
1243−123......... 12 46 00.69 −12 36 19.9 0.406 305.4 S SCR 1246-1236 = · · · = · · ·
1316−215......... 13 19 24.72 −21 47 55.0 0.467 179.2 S NLTT 33669 = LP 854-50 = WT 2034
1436−781......... 14 42 51.54 −78 23 53.6 0.409 272.0 S NLTT 38003 = LP 40-109 = LTT 5814
1452−310......... 14 55 23.47 −31 17 06.4 0.199 174.2 P LEHPM 2-4029 = · · · = · · ·
1647−327......... 16 50 44.32 −32 49 23.2 0.526 193.8 L LHS 3245 = NLTT 43628 = LP 919-1
1742−722......... 17 48 31.21 −72 17 18.5 0.294 228.2 P LEHPM 2-1166 = · · · = · · ·
1946−273......... 19 49 19.78 −27 12 25.7 0.213 162.0 L NLTT 48270 = LP 925-53 = · · ·
2008−600......... 20 12 31.75 −59 56 51.5 1.440 165.6 S SCR 2012-5956 = · · · = · · ·
2008−799......... 20 16 49.66 −79 45 53.0 0.434 128.4 S SCR 2016-7945 = · · · = · · ·
2035−369......... 20 38 41.42 −36 49 13.5 0.230 104.0 L NLTT 49589 = L 495-42 = LEHPM 2-3290
2103−397......... 21 06 32.01 −39 35 56.7 0.266 151.7 P LEHPM 2-1571 = · · · = · · ·
2138−332......... 21 41 57.56 −33 00 29.8 0.210 228.5 P NLTT 51844 = L 570-26 = LEHPM 2-3327
2157−574......... 22 00 45.37 −57 11 23.4 0.233 252.0 P LEHPM 1-4327 = · · · = · · ·
2218−416......... 22 21 25.37 −41 25 27.0 0.210 143.4 P LEHPM 1-4598 = · · · = · · ·
2231−387......... 22 33 54.47 −38 32 36.9 0.370 220.5 P NLTT 54169 = LP 1033-28 = LEHPM 1-4859
Known White Dwarfs without a Trigonometric Parallax Estimated to be Within 25 pc
0141−675 ........ 01 43 00.98 −67 18 30.3 1.048 197.8 L LHS 145 = NLTT 5777 = L 88-59
0806−661 ........ 08 06 53.76 −66 18 16.6 0.454 131.4 S NLTT 19008 = L 97-3 = · · ·
1009−184 ........ 10 12 01.88 −18 43 33.2 0.519 268.2 S WT 1759 = LEHPM 2-220 = · · ·
1036−204 ........ 10 38 55.57 −20 40 56.7 0.628 330.3 L LHS 2293 = NLTT 24944 = LP 790-29
1202−232 ........ 12 05 26.66 −23 33 12.1 0.227 002.0 L NLTT 29555 = LP 852-7 = LEHPM 2-1894
1315−781 ........ 13 19 25.63 −78 23 28.3 0.477 139.2 S NLTT 33551 = L 40-116 = · · ·
1339−340 ........ 13 42 02.88 −34 15 19.4 2.547 296.7 Le PM J13420-3415 = · · · = · · ·
1756+143 ........ 17 58 22.90 +14 17 37.8 1.014 235.4 Le LSR 1758+1417 = · · · = · · ·
1814+134 ........ 18 17 06.48 +13 28 25.0 1.207 201.5 Le LSR 1817+1328 = · · · = · · ·
2040−392 ........ 20 43 49.21 −39 03 18.0 0.306 179.0 L NLTT 49752 = L 495-82 = · · ·
2211−392 ........ 22 14 34.75 −38 59 07.3 1.056 110.1 O WD J2214-390 = LEHPM 1-4466 = · · ·
2226−754A........ 22 30 40.00 −75 13 55.3 1.868 167.5 S SSSPM J2231-7514 = · · · = · · ·
2226−754B........ 22 30 33.55 −75 15 24.2 1.868 167.5 S SSSPM J2231-7515 = · · · = · · ·
Known White Dwarfs without a Trigonometric Parallax Estimated to be Beyond 25 pc
0024−556......... 00 26 40.69 −55 24 44.1 0.580 211.8 L LHS 1076 = NLTT 1415 = L 170-27
0150+256......... 01 52 51.93 +25 53 40.7 0.220 076.0 L NLTT 6275 = G 94-21 = · · ·
0255−705......... 02 56 17.22 −70 22 10.8 0.682 097.9 L LHS 1474 = NLTT 9485 = L 54-5
0442−304......... 04 44 29.38 −30 21 14.2 0.196 199.5 P NLTT 13882 = LP 891-65 = HE 0442-3027
0928−713......... 09 29 07.97 −71 33 58.8 0.439 320.2 S NLTT 21957 = L 64-40 = · · ·
1143−013......... 11 46 25.77 −01 36 36.8 0.563 140.2 S LHS 2455 = NLTT 28493 = · · ·
– 21 –
Table 3—Continued
WD Name RA Dec PM PA Ref Alternate Names
(J2000.0) (J2000.0) (arcsec yr−1) (deg)
1237−230......... 12 40 24.18 −23 17 43.8 1.102 219.9 L LHS 339 = NLTT 31473 = LP 853-15
1314−153......... 13 16 43.59 −15 35 58.3 0.708 196.7 L LHS 2712 = NLTT 33503 = LP 737-47
1418−088......... 14 20 54.93 −09 05 08.7 0.480 266.8 S LHS 5270 = NLTT 37026 = · · ·
1447−190......... 14 50 11.93 −19 14 08.7 0.253 285.4 P NLTT 38499 = LP 801-14 = LEHPM 2-1835
1607−250......... 16 10 50.21 −25 13 16.0 0.209 314.0 L NLTT 42153 = LP 861-31 = · · ·
References. — (L) Luyten 1979a,b, (Le) Lépine et al. 2003, Lépine et al. 2005, (O) Oppenheimer et al. 2001, (P) Pokorny et al.
2004, (S) Subasavage et al. 2005a,b, this work
– 22 –
Fig. 1.— Reduced proper motion diagram used to select WD candidates for spectroscopic
follow-up. Plotted are the new high proper motion objects from Subasavage et al. (2005a,b).
The line is a somewhat arbitrary boundary between the WDs (below) and the subdwarfs
(just above). Main sequence dwarfs fall above and to the right of the subdwarfs, although
there is significant overlap. Asterisks indicate the 33 new WDs reported here. Three dots
in the WD region are deferred to a future paper. The point labeled “sd” is a confirmed
subdwarf contaminant of the WD sample.
Fig. 2.— Spectral plots of the hot (Teff ≥ 10000 K) DA WDs from the new sample, plotted
in descending Teff as derived from the SED fits to the photometry. Note that some of the
flux calibrations are not perfect, in particular, at the blue end.
Fig. 3.— Spectral plots of cool (Teff < 10000 K) DA WDs from the new sample, plotted in
descending Teff as derived from the SED fits to the photometry. Note that some of the flux
calibrations are not perfect, in particular, at the blue end.
Fig. 4.— Spectral plots of the four featureless DC white dwarfs from the new sample, plotted
in descending Teff as derived from the SED fits to the photometry. Note that some of the
flux calibrations are not perfect, in particular, at the blue end.
Fig. 5.— Comparison plot of the values of Teff derived from photometric SED fitting vs
those derived from spectral fitting for 25 of the DA WDs in the new sample. The solid line
represents equal temperatures. The elevated point, 0310−624, is discussed in § 4.2.
Fig. 6.— Spectral plot of WD 0121−429. The inset plot displays the spectrum (light line)
in the Hα region to which a magnetic fit (heavy line), as outlined in Bergeron et al. (1992a),
was performed using the Teff obtained from the SED fit to the photometry. The resulting
magnetic parameters are listed below the fit.
Fig. 7.— Spectral plot of WD 0622−329. The inset plot displays the spectrum (light line)
in the region to which the model (heavy line) was fit assuming the spectrum is a convolution
of a DB component and a slightly cooler DA component. Best fit physical parameters are
listed below the fit for each component.
Fig. 8.— (top panel) Spectral plot of WD 0840−136. The DZ model failed to reproduce
the spectrum presumably because this object is cooler than Teff ∼ 5000 K where additional
pressure effects, not included in the model, become important. (bottom panel) Spectral plot
of WD 2138−332. The inset plot displays the spectrum (light line) in the region to which
the model (heavy line) was fit.
– 23 –
Fig. 9.— Spectral plot of WD 1149−272. The inset plot displays the spectrum (light line)
in the region to which the model (heavy line) was fit.
Fig. 10.— Spectral energy distribution plot of WD 2008−600 with the distance constrained
by the trigonometric distance of 17.1 ± 0.4 pc. Best fit physical parameters are listed below
the fit. Points are fit values; error bars are derived from the uncertainties in the magnitudes
and the parallax.
– 24 –
– 25 –
– 26 –
– 27 –
– 28 –
– 29 –
– 30 –
– 31 –
– 32 –
– 33 –
	Introduction
	Candidate Selection
	Data and Observations
	Astrometry and Nomenclature
	Spectroscopy
	Photometry
	Analysis
	Modeling of Physical Parameters
	Comments on Individual Systems
	Discussion
	Acknowledgments
	Appendix

  We present spectra for 33 previously unclassified white dwarf systems
brighter than V = 17 primarily in the southern hemisphere. Of these new
systems, 26 are DA, 4 are DC, 2 are DZ, and 1 is DQ. We suspect three of these
systems are unresolved double degenerates. We obtained VRI photometry for these
33 objects as well as for 23 known white dwarf systems without trigonometric
parallaxes, also primarily in the southern hemisphere. For the 56 objects, we
converted the photometry values to fluxes and fit them to a spectral energy
distribution using the spectroscopy to determine which model to use (i.e. pure
hydrogen, pure helium, or metal-rich helium), resulting in estimates of
effective temperature and distance. Eight of the new and 12 known systems are
estimated to be within the NStars and Catalogue of Nearby Stars (CNS) horizons
of 25 pc, constituting a potential 18% increase in the nearby white dwarf
sample. Trigonometric parallax determinations are underway via CTIOPI for these
20 systems.
  One of the DCs is cool so that it displays absorption in the near infrared.
Using the distance determined via trigonometric parallax, we are able to
constrain the model-dependent physical parameters and find that this object is
most likely a mixed H/He atmosphere white dwarf similar to other cool white
dwarfs identified in recent years with significant absorption in the infrared
due to collision-induced absorptions by molecular hydrogen.

Introduction
The study of white dwarfs (WDs) provides insight to understanding WD formation
rates, evolution, and space density. Cool WDs, in particular, provide limits on the age of the
Galactic disk and could represent some unknown fraction of the Galactic halo dark matter.
Individually, nearby WDs are excellent candidates for astrometric planetary searches because
the astrometric signature is greater than for an identical WD system more distant. As a
population, a complete volume limited sample is necessary to provide unbiased statistics;
however, their intrinsic faintness has allowed some to escape detection.
Of the 18 WDs with trigonometric parallaxes placing them within 10 pc of the Sun (the
RECONS sample), all but one have proper motions greater than 1.′′0 yr−1 (94%). By com-
parison, of the 230 main sequence systems (as of 01 January 2007) in the RECONS sample,
50% have proper motions greater than 1.′′0 yr−1. We have begun an effort to reduce this
apparent selection bias against slower-moving WDs to complete the census of nearby WDs.
This effort includes spectroscopic, photometric, and astrometric initiatives to characterize
newly discovered as well as known WDs without trigonometric parallaxes. Utilizing the Su-
perCOSMOS Sky Survey (SSS) for plate magnitude and proper motion information coupled
with data from other recently published proper motion surveys (primarily in the southern
hemisphere), we have identified relatively bright WD candidates via reduced proper motion
diagrams.
– 3 –
In this paper, we present spectra for 33 newly discovered WD systems brighter than
V = 17.0. Once an object is spectroscopically confirmed to be a WD (in this paper for
the first time or elsewhere in the literature), we obtain CCD photometry to derive Teff and
estimate its distance using a spectral energy distribution (SED) fit and a model atmosphere
analysis. If an object’s distance estimate is within the NStars (Henry et al. 2003) and CNS
(Gliese & Jahreiß 1991) horizons of 25 pc, it is then added to CTIOPI (Cerro Tololo Inter-
American Observatory Parallax Investigation) to determine its true distance (e.g. Jao et al.
2005, Henry et al. 2006).
2. Candidate Selection
We used recent high proper motion (HPM) surveys (Pokorny et al. 2004; Subasavage et al.
2005a,b; Finch et al. 2007) in the southern hemisphere for this work because our long-term
astrometric observing program CTIOPI, is based in Chile. To select good WD candidates
for spectroscopic observations, plate magnitudes via SSS and 2MASS JHKS are extracted
for HPM objects. Each object’s (R59F − J) color and reduced proper motion (RPM) are
then plotted. RPM correlates proper motion with proximity, which is certainly not always
true; however, it is effective at separating WDs from subdwarfs and main sequence stars.
Figure 1 displays an RPM diagram for the 33 new WDs presented here. To serve as examples
for the locations of subdwarfs and main sequence stars, recent HPM discoveries from the
SuperCOSMOS-RECONS (SCR) proper motion survey are also plotted (Subasavage et al.
2005a,b). The solid line represents a somewhat arbitrary cutoff separating subdwarfs and
WDs. Targets are selected from the region below the solid line. Note there are four stars
below this line that are not represented with asterisks. Three have recently been spectro-
scopically confirmed as WDs (Subasavage et al., in preparation) and one as a subdwarf (SCR
1227−4541, denoted by “sd”) that fell just below the line at (R59F − J) = 1.4 and HR59F =
19.8 (Subasavage et al. 2005b).
Completeness limits (S/N > 10) for 2MASS are J = 15.8, H = 15.1, and KS = 14.3
for uncontaminated point sources (Skrutskie et al. 2006). The use of J provides a more
reliable RPM diagram color for objects more than a magnitude fainter than the KS limit,
which is particularly important for the WDs (with (J − KS) < 0.4) discussed here. Only
objects bright enough to have 2MASS magnitudes are included in Figure 1. Consequently,
all WD candidates are brighter than V ∼ 17, and are therefore likely to be nearby. Objects
that fall in the WD region of the RPM diagram were cross-referenced with SIMBAD and
– 4 –
McCook & Sion (1999)1 to determine those that were previously classified as WDs. The
remainder were targeted for spectroscopic confirmation.
The remaining 33 candidates comprise the “new sample” whose spectra are presented
in this work, while the “known sample” constitutes the 23 previously identified WD systems
without trigonometric parallaxes for which we have complete V RIJHKS data.
3. Data and Observations
3.1. Astrometry and Nomenclature
The traditional naming convention for WDs uses the object’s epoch 1950 equinox 1950
coordinates. Coordinates for the new sample were extracted from 2MASS along with the
Julian date of observation. These coordinates were adjusted to account for proper motion
from the epoch of 2MASS observation to epoch 2000 (hence epoch 2000 equinox 2000).
The coordinates were then transformed to equinox 1950 coordinates using the IRAF proce-
dure precess. Finally, the coordinates were again adjusted (opposite the direction of proper
motion) to obtain epoch 1950 equinox 1950 coordinates.
Proper motions were taken from various proper motion surveys in addition to unpub-
lished values obtained via the SCR proper motion survey while recovering previously known
HPM objects. Appendix A contains the proper motions used for coordinate sliding as well
as J2000 coordinates and alternate names.
3.2. Spectroscopy
Spectroscopic observations were taken on five separate observing runs in 2003 Octo-
ber and December, 2004 March and September, and 2006 May at the Cerro Tololo Inter-
American Observatory (CTIO) 1.5m telescope as part of the SMARTS Consortium. The
Ritchey-Chrétien Spectrograph and Loral 1200×800 CCD detector were used with grating
09, providing 8.6 Å resolution and wavelength coverage from 3500 to 6900 Å. Observations
consisted of two exposures (typically 20 - 30 minutes each) to permit cosmic ray rejection,
followed by a comparison HeAr lamp exposure to calibrate wavelength for each object. Bias
subtraction, dome/sky flat-fielding, and extraction of spectra were performed using standard
IRAF packages.
1The current web based catalog can be found at http://heasarc.nasa.gov/W3Browse/all/mcksion.html
http://heasarc.nasa.gov/W3Browse/all/mcksion.html
– 5 –
A slit width of 2′′ was used for the 2003 and 2004 observing runs. Some of these data have
flux calibration problems because the slit was not rotated to be aligned along the direction
of atmospheric refraction. In conjunction with telescope “jitter”, light was sometimes lost
preferentially at the red end or the blue end for these data.
A slit width of 6′′, used for the 2006 May run, eliminated most of the flux calibration
problems even though the slit was not rotated. All observations were taken at an airmass
of less than 2.0. Within our wavelength window, the maximum atmospheric differential
refraction is less than 3′′ (Filippenko 1982). A test was performed to verify that no resolution
was lost by taking spectra of a F dwarf with sharp absorption lines from slit widths of 2′′ to
10′′ in 2′′ increments. Indeed, no resolution was lost.
Spectra for the new DA WDs with Teff ≥ 10000 K are plotted in Figure 2 while spectra
for the new DA WDs with Teff < 10000 K are plotted in Figure 3. Featureless DC spectra
are plotted in Figure 4. Spectral plots as well as model fits for unusual objects are described
in § 4.2.
3.3. Photometry
Optical V RI (Johnson V , Kron-Cousins RI) for the new and known samples was ob-
tained using the CTIO 0.9 m telescope during several observing runs from 2003 through
2006 as part of the Small and Moderate Aperture Research Telescope System (SMARTS)
Consortium. The 2048×2046 Tektronix CCD camera was used with the Tek 2 V RI filter
set2. Standard stars from Graham (1982), Bessel (1990), and Landolt (1992) were observed
each night through a range of airmasses to calibrate fluxes to the Johnson-Kron-Cousins
system and to calculate extinction corrections.
Bias subtraction and dome flat-fielding (using calibration frames taken at the beginning
of each night) were performed using standard IRAF packages. When possible, an aperture
14′′ in diameter was used to determine the stellar flux, which is consistent with the aperture
used by Landolt (1992) for the standard stars. If cosmic rays fell within this aperture, they
were removed before flux extraction. In cases of crowded fields, aperture corrections were
applied and ranged from 4′′ to 12′′ in diameter using the largest aperture possible without
including contamination from neighboring sources. Uncertainties in the optical photometry
were derived by estimating the internal night-to-night variations as well as the external errors
(i.e. fits to the standard stars). A complete discussion of the error analysis can be found in
2The central wavelengths for V , R, and I are 5475, 6425, and 8075Å respectively.
– 6 –
Henry et al. (2004). We adopt a total error of ±0.03 mag in each band. The final optical
magnitudes are listed in Table 1 as well as the number of nights each object was observed.
Infrared JHKS magnitudes and errors were extracted via Aladin from 2MASS and are
also listed in Table 1. JHKS magnitude errors are, in most cases, significantly larger than for
V RI, and the errors listed give a measure of the total photometric uncertainty (i.e. include
both global and systematic components). In cases when the magnitude error is null, the star
is near the magnitude limit of 2MASS and the photometry is not reliable.
4. Analysis
4.1. Modeling of Physical Parameters
The pure hydrogen, pure helium, and mixed hydrogen and helium model atmospheres
used to model the WDs are described at length in Bergeron et al. (2001) and references
therein, while the helium-rich models appropriate for DQ and DZ stars are described in
Dufour et al. (2005, 2007), respectively. The atmospheric parameters for each star are ob-
tained by converting the optical V RI and infrared JHKS magnitudes into observed fluxes,
and by comparing the resulting SEDs with those predicted from our model atmosphere cal-
culations. The first step is accomplished by transforming the magnitudes into average stellar
fluxes fm
received at Earth using the calibration of Holberg et al. (2006) for photon count-
ing devices. The observed and model fluxes, which depend on Teff , log g, and atmospheric
composition, are related by the equation
= 4π (R/D)2 Hm
, (1)
where R/D is the ratio of the radius of the star to its distance from Earth, and Hm
the Eddington flux, properly averaged over the corresponding filter bandpass. Our fitting
technique relies on the nonlinear least-squares method of Levenberg-Marquardt (Press et al.
1992), which is based on a steepest descent method. The value of χ2 is taken as the sum over
all bandpasses of the difference between both sides of eq. (1), weighted by the corresponding
photometric uncertainties. We consider only Teff and the solid angle to be free parameters,
and the uncertainties of both parameters are obtained directly from the covariance matrix
of the fit. In this study, we simply assume a value of log g = 8.0 for each star.
As discussed in Bergeron et al. (1997, 2001), the main atmospheric constituent — hy-
drogen or helium — is determined by comparing the fits obtained with both compositions,
or by the presence of Hα in the optical spectra. For DQ and DZ stars, we rely on the
– 7 –
procedure outlined in Dufour et al. (2005, 2007), respectively: we obtain a first estimate
of the atmospheric parameters by fitting the energy distribution with an assumed value of
the metal abundances. We then fit the optical spectrum to measure the metal abundances,
and use these values to improve our atmospheric parameters from the energy distribution.
This procedure is iterated until a self-consistent photometric and spectroscopic solution is
achieved.
The derived values for Teff for each object are listed in Table 1. Also listed are the
spectral types for each object determined based on their spectral features. The DAs have
been assigned a half-integer temperature index as defined by McCook & Sion (1999), where
the temperature index equals 50,400/Teff. As an external check, we compare in Figure 5
the photometric effective temperatures for the DA stars in Table 1 with those obtained by
fitting the observed Balmer line profiles (Figs. 2 and 3) using the spectroscopic technique
developed by Bergeron et al. (1992b), and recently improved by Liebert et al. (2003). Our
grid of pure hydrogen, NLTE, and convective model atmospheres is also described in Liebert
et al. The uncertainties of the spectroscopic technique are typically of 0.038 dex in log g
and 1.2% in Teff according to that study. We adopt a slightly larger uncertainty of 1.5%
in Teff (Spec) because of the problematic flux calibrations of the pre−2006 data (see § 3.2).
The agreement shown in Figure 5 is excellent, except perhaps at high temperatures where
the photometric determinations become more uncertain. It is possible that the significantly
elevated point in Figure 5, WD 0310−624 (labeled), is an unresolved double degenerate (see
§ 4.2). We refrain here from using the log g determinations in our analysis because these
are available only for the DA stars in our sample, and also because the spectra are not flux
calibrated accurately enough for that purpose.
Once the effective temperature and the atmospheric composition are determined, we
calculate the absolute visual magnitude of each star by combining the new calibration of
Holberg et al. (2006) with evolutionary models similar to those described in Fontaine et al.
(2001) but with C/O cores, q(He) ≡ logMHe/M⋆ = 10
−2 and q(H) = 10−4 (representative
of hydrogen-atmosphere WDs), and q(He) = 10−2 and q(H) = 10−10 (representative of
helium-atmosphere WDs)3. By combining the absolute visual magnitude with the Johnson
V magnitude, we derive a first estimate of the distance of each star (reported in Table 1).
Errors on the distance estimates incorporate the errors of the photometry values as well as
an error of 0.25 dex in log g, which is the measured dispersion of the observed distribution
using spectroscopic determinations (see Figure 9 of Bergeron et al. 1992b).
Of the 33 new systems presented here, 5 have distance estimates within 25 pc. Four
3see http://www.astro.umontreal.ca/˜bergeron/CoolingModels/
http://www.astro.umontreal.ca/~bergeron/CoolingModels/
– 8 –
more systems require additional attention because distance estimates are derived via other
means. Three of these are likely within 25 pc. All four are further discussed in the next
section. In total, 20 WD systems (8 new and 12 known) are estimated (or determined) to be
within 25 pc and one additional common proper motion binary system possibly lies within
25 pc.
4.2. Comments on Individual Systems
Here we address unusual and interesting objects.
WD 0121−429 is a DA WD that exhibits Zeeman splitting of Hα and Hβ, thereby
making its formal classification DAH. The SED fit to the photometry is superb, yielding
a Teff of 6,369 ± 137 K. When we compare the strength of the absorption line trio with
that predicted using the Teff from the SED fit, the depth of the absorption appears too
shallow. Using the magnetic line fitting procedure outlined in Bergeron et al. (1992a), we
must include a 50% dilution factor to match the observed central line of Hα. In light of
this, we utilized the trigonometric parallax distance determined via CTIOPI of 17.7 ± 0.7
pc (Subasavage et al., in preparation) to further constrain this system. The resulting SED
fit, with distance (hence luminosity) as a constraint rather than a variable, implies a mass of
0.43 ± 0.03 M⊙. Given the age of our Galaxy, the lowest mass WD that could have formed
is ∼0.47 M⊙ (Iben & Renzini 1984). It is extremely unlikely that this WD formed through
single star evolution. The most likely scenario is that this is a double degenerate binary
with a magnetic DA component and a featureless DC component (necessary to dilute the
absorption at Hα), similar to G62-46 (Bergeron et al. 1993) and LHS 2273 (see Figure 33
of Bergeron et al. 1997). If this interpretation is correct, any number of component masses
and luminosities can reproduce the SED fit.
The spectrum and corresponding magnetic fit to the Hα lines (including the dilution)
is shown in Figure 6. The viewing angle, i = 65◦, is defined as the angle between the dipole
axis and the line of sight (i = 0 corresponds to a pole-on view). The best fit produces a
dipole field strength, Bd = 9.5 MG, and a dipole offset, az = 0.06 (in units of stellar radius).
The positive value of az implies that the offset is toward the observer. Only Bd is moderately
constrained, both i and az can vary significantly yet still produce a reasonable fit to the data
(Bergeron et al. 1992a).
WD 0310−624 is a DAWD that is one of the hottest in the new sample. Because of it’s
elevation significantly above the equal temperature line (solid) in Figure 5, it is possible that
it is an unresolved double degenerate with very different component effective temperatures.
– 9 –
In fact, this method has been used to identify unresolved double degenerate candidates (i.e.
Bergeron et al. 2001).
WD 0511−415 is a DA WD (spectrum is plotted in Figure 2) whose spectral fit
produces a Teff = 10,813 ± 219 K and a log g = 8.21 ± 0.10 using the spectral fitting
procedure of Liebert et al. (2003). This object lies near the red edge of the ZZ Ceti instability
strip as defined by Gianninas et al. (2006). If variable, this object would help to constrain
the cool edge of the instability strip in Teff , log g parameter space. Follow-up high speed
photometry is necessary to confirm variability.
WD 0622−329 is a DAB WD displaying the Balmer lines as well as weaker He I at
4472 and 5876 Å. The spectrum, shown in Figure 7, is reproduced best with a model having
Teff ∼43,700 K. However, the predicted He II absorption line at 4686 Å for a WD of this
Teff is not present in the spectrum. In contrast, the SED fit to the photometry implies a
Teff of ∼10,500 K (using either pure H or pure He models). Because the Teff values are
vastly discrepant, we explored the possibility that this spectrum is not characterized by a
single temperature. We modeled the spectrum assuming the object was an unresolved double
degenerate. The best fit implies one component is a DB with Teff = 14,170 ± 1,228 K and
the other component is a DA with Teff = 9,640 ± 303 K, similar to the unresolved DA +
DB degenerate binary PG 1115+166 analyzed by Bergeron & Liebert (2002). One can see
from Figure 7 that the spectrum is well modeled under this assumption. We conclude this
object is likely a distant (well beyond 25 pc) unresolved double degenerate.
WD 0840−136 is a DZ WD whose spectrum shows both Ca II (H & K) and Ca I (4226
Å) lines as shown in Figure 8. Fits to the photometric data for different atmospheric com-
positions indicate temperatures of about 4800-5000 K. However, fits to the optical spectrum
using the models of Dufour et al. (2007) cannot reproduce simultaneously all three calium
lines. This problem is similar to that encountered by Dufour et al. (2007) where the atmo-
spheric parameters for the coolest DZ WDs were considered uncertain because of possible
high atmospheric pressure effects. We utilize a photometric relation relevant for WDs of any
atmospheric composition, which links MV to (V −I) (Salim et al. 2004) to obtain a distance
estimate of 19.3 ± 3.9 pc.
WD 1054−226 was observed spectroscopically as part of the Edinburgh-Cape (EC)
blue object survey and assigned a spectral type of sdB+ (Kilkenny et al. 1997). As is evident
in Figure 3, the spectrum of this object is the noisiest of all the spectra presented here and
perhaps a bit ambiguous. As an additional check, this object was recently observed using
the ESO 3.6 m telescope and has been confirmed to be a cool DA WD (Bergeron, private
communication).
– 10 –
WD 1105−340 is a DA WD (spectrum is plotted in Figure 2) with a common proper
motion companion with separation of 30.′′6 at position angle 107.1◦. The companion’s spectral
type is M4Ve with VJ = 15.04, RKC = 13.68, IKC = 11.96, J = 10.26, H = 9.70, and KS
= 9.41. In addition to the SED derived distance estimate for the WD, we utilize the main
sequence distance relations of Henry et al. (2004) to estimate a distance to the red dwarf
companion. We obtain a distance estimate of 19.1 ± 3.0 pc for the companion leaving
open the possibility that this system may lie just within 25 pc. A trigonometric parallax
determination is currently underway for confirmation.
WD 1149−272 is the only DQ WD discovered in the new sample. This object was
observed spectroscopically as part of the Edinburgh-Cape (EC) blue object survey for which
no features deeper than 5% were detected and was labeled a possible DC (Kilkenny et al.
1997). It is identified as having weak C2 swan band absorption at 4737 and 5165 Å and is
otherwise featureless. The DQ model reproduces the spectrum reliably and is overplotted in
Figure 9. This object is characterized as having Teff = 6188 ± 194 K and a log (C/He) =
−7.20 ± 0.16.
WD 2008−600 is a DC WD (spectrum is plotted in Figure 4) that is flux deficient in
the near infrared, as indicated by the 2MASS magnitudes. The SED fit to the photometry
is a poor match to either the pure hydrogen or the pure helium models. A pure hydrogen
model provides a slightly better match than a pure helium model, and yields a Teff of ∼3100
K, thereby placing it in the relatively small sample of ultracool WDs. In order to discern the
true nature of this object, we have constrained the model using the distance obtained from
the CTIOPI trigonometric parallax of 17.1 ± 0.4 pc (Subasavage et al., in preparation). This
object is then best modeled as having mostly helium with trace amounts of hydrogen (log
(He/H) = 2.61) in its atmosphere and has a Teff = 5078 ± 221 K (see Figure 10). A mixed
hydrogen and helium composition is required to produce sufficient absorption in the infrared
as a result of the collision-induced absorption by molecular hydrogen due to collisions with
helium. Such mixed atmospheric compositions have also been invoked to explain the infrared
flux deficiency in LHS 1126 (Bergeron et al. 1994) as well as SDSS 1337+00 and LHS 3250
(Bergeron & Leggett 2002). While WD 2008−600 is likely not an ultracool WD, it is one of
the brightest and nearest cool WDs known. Because the 2MASS magnitudes are not very
reliable, we intend to obtain additional near-infrared photometry to better constrain the fit.
WD 2138−332 is a DZ WD for which a calcium rich model reproduces the spectrum
reliably. The spectrum and the overplotted fit are shown in the bottom panel of Figure
8. Clearly evident in the spectrum are the strong Ca II absorption at 3933 and 3968 Å. A
weaker Ca I line is seen at 4226Å. Also seen are Mg I absorption lines at 3829, 3832, and
3838 Å (blended) as well as Mg I at 5167, 5173, and 5184 Å (also blended). Several weak Fe I
– 11 –
lines from 4000Å to 4500Å and again from 5200Å to 5500Å are also present. The divergence
of the spectrum from the fit toward the red end is likely due to an imperfect flux calibration
of the spectrum. This object is characterized as having Teff = 7188 ± 291 K and a log
(Ca/He) = −8.64 ± 0.16. The metallicity ratios are, at first, assumed to be solar (as defined
by Grevesse & Sauval 1998) and, in this case, the quality of the fit was sufficient without
deviation. The corresponding log (Mg/He) = −7.42 ± 0.16 and log (Fe/He) = −7.50 ± 0.16
for this object.
WD 2157−574 is a DAWD (spectrum is plotted in Figure 3) unique to the new sample
in that it displays weak Ca II absorption at 3933 and 3968 Å (H and K) thereby making its
formal classification a DAZ. Possible scenarios that enrich the atmospheres of DAZs include
accretion via (1) debris disks, (2) ISM, and (3) cometary impacts (see Kilic et al. 2006 and
references therein). The 2MASS KS magnitude is near the faint limit and is unreliable, but
even considering the J and H magnitudes, there appears to be no appreciable near-infrared
excess. While this may tentatively rule out the possibility of a debris disk, this object would
be an excellent candidate for far-infrared spaced-based studies to ascertain the origin of the
enrichment.
5. Discussion
WDs represent the end state for stars less massive than ∼8 M⊙ and are therefore rel-
atively numerous. Because of their intrinsic faintness, only the nearby WD population can
be easily characterized and provides the benchmark upon which WD stellar astrophysics is
based. It is clear from this work and others (e.g. Holberg et al. 2002; Kawka & Vennes 2006)
that the WD sample is complete, at best, to only 13 pc. Spectroscopic confirmation of new
WDs as well as trigonometric parallax determinations for both new and known WDs will
lead to a more complete sample and will push the boundary of completeness outward. We
estimate that 8 new WDs and an additional 12 known WDs without trigonometric parallaxes
are nearer than 25 pc, including one within 10 pc (WD 0141−675). Parallax measurements
via CTIOPI are underway for these 20 objects to confirm proximity. This total of 20 WDs
within 25 pc constitutes an 18% increase to the 109 WDs with trigonometric parallaxes ≥
40 mas.
Evaluating the proper motions of the new and known samples within 25 pc indicates
that almost double the number of systems have been found with µ < 1.′′0 yr−1 than with
µ ≥ 1.′′0 yr−1 (13 vs 7, see Table 2). The only WD estimated to be within 10 pc has µ >
1.′′0 yr−1, although WD 1202−232 is estimated to be 10.2 ± 1.7 pc and it’s proper motion is
small (µ = 0.′′227 yr−1).
– 12 –
Because this effort focuses mainly on the southern hemisphere, it is likely that there
is a significant fraction of nearby WDs in the northern hemisphere that have also gone
undetected. With the recent release of the LSPM-North Catalog (Lépine & Shara 2005),
these objects are identifiable by employing the same techniques used in this work. The
challenge is the need for a large scale parallax survey focusing on WDs to confirm proximity.
Since the HIPPARCOS mission, only six WD trigonometric parallaxes have been published
(Hambly et al. 1999; Smart et al. 2003), and of those, only two are within 25 pc. The USNO
parallax program is in the process of publishing trigonometric parallaxes for ∼130 WDs,
mostly in the northern hemisphere, although proximity was not a primary motivation for
target selection (Dahn, private communication).
In addition to further completing the nearby WD census, the wealth of observational
data available from this effort provides reliable constraints on their physical parameters
(i.e. Teff , log g, mass, and radius). Unusual objects are then revealed, such as those dis-
cussed in § 4.2. In particular, trigonometric parallaxes help identify WDs that are overlu-
minous, as is the case for WD 0121−429. This object, and others similar to it, are excellent
candidates to provide insight into binary evolution. If they can be resolved using high res-
olution astrometric techniques (i.e. speckle, adaptive optics, or interferometry via Hubble
Space Telescope’s Fine Guidance Sensors), they may provide astrometric masses, which are
fundamental calibrators for stellar structure theory and for the reliability of the theoretical
WD mass-radius and initial-to-final-mass relationships. To date, only four WD astrometric
masses are known to better than ∼ 5% (Provencal et al. 1998).
One avenue that is completely unexplored to date is a careful high resolution search
for planets around WDs. Theory dictates that the Sun will become a WD, and when it
does, the outer planets will remain in orbit (not without transformations of their own, of
course). In this scenario, the Sun will have lost more than half of its mass, thereby amplifying
the signature induced by the planets. Presumably, this has already occurred in the Milky
Way and systems such as these merely await detection. Because of the faintness and spectral
signatures of WDs (i.e. few, if any, broad absorption lines), current radial velocity techniques
are inadequate for planet detection, leaving astrometric techniques as the only viable option.
For a given system, the astrometric signature is inversely related to distance (i.e. the nearer
the system, the larger the astrometric signature). This effort aims to provide a complete
census of nearby WDs that can be probed for these astrometric signatures using future
astrometric efforts.
– 13 –
6. Acknowledgments
The RECONS team at Georgia State University wishes to thank the NSF (grant AST
05-07711), NASA’s Space Interferometry Mission, and GSU for their continued support of
our study of nearby stars. We also thank the continuing support of the members of the
SMARTS consortium, who enable the operations of the small telescopes at CTIO where all
of the data in this work were collected. J. P. S. is indebted to Wei-Chun Jao for the use of his
photometry reduction pipeline. P. B. is a Cottrell Scholar of Research Corporation and would
like to thank the NSERC Canada for its support. N. C. H. would like to thank colleagues in
the Wide Field Astronomy Unit at Edinburgh for their efforts contributing to the existence
of the SSS; particular thanks go to Mike Read, Sue Tritton, and Harvey MacGillivray. This
work has made use of the SIMBAD, VizieR, and Aladin databases, operated at the CDS in
Strasbourg, France. We have also used data products from the Two Micron All Sky Survey,
which is a joint project of the University of Massachusetts and the Infrared Processing and
Analysis Center, funded by NASA and NSF.
A. Appendix
In order to ensure correct cross-referencing of names for the new and known WD systems
presented here, Table 3 lists additional names found in the literature. Objects for which there
is an NLTT designation will also have the corresponding L or LP designations found in the
NLTT catalog. This is necessary because the NLTT designations were not published in the
original catalog, but rather are the record numbers in the electronic version of the catalog
and have been adopted out of necessity.
REFERENCES
Bergeron, P., Ruiz, M.-T., & Leggett, S. K. 1992, ApJ, 400, 315
Bergeron, P., Saffer, R. A., & Liebert, J. 1992, ApJ, 394, 228
Bergeron, P., Ruiz, M.-T., & Leggett, S. K. 1993, ApJ, 407, 733
Bergeron, P., Ruiz, M.-T., Leggett, S. K., Saumon, D., & Wesemael, F. 1994, ApJ, 423, 456
Bergeron, P., Ruiz, M. T., & Leggett, S. K. 1997, ApJS, 108, 339
Bergeron, P., Leggett, S. K., & Ruiz, M. T. 2001, ApJS, 133, 413
– 14 –
Bergeron, P., & Leggett, S. K. 2002, ApJ, 580, 1070
Bergeron, P., & Liebert, J. 2002, ApJ, 566, 1091
Bessel, M. S. 1990, A&AS, 83, 357
Carpenter, J. M. 2001, AJ, 121, 2851
Dufour, P., Bergeron, P., & Fontaine, G. 2005, ApJ, 627, 404
Dufour, P., Bergeron, P., Liebert, J., Harris, H. C., Knapp, G. R., Anderson, S. F., Hall,
P. B., Strauss, M. A., Collinge, M. J., & Edwards, M. C. 2007, submitted to ApJ
Filippenko, A. V. 1982, PASP, 94, 715
Finch, C. T., Henry, T. J., Subasavage, J. P., Jao, W.-C., Hambly, N. C. 2007, AJ, submitted
Fontaine, G., Brassard, P., & Bergeron, P. 2001, PASP, 113, 409
Gianninas, A., Bergeron, P., & Fontaine, G. 2006, AJ, 132, 831
Gliese, W., & Jahreiß, H. 1991, On: The Astronomical Data Center CD-ROM: Selected As-
tronomical Catalogs, Vol. I; L.E. Brotzmann, S.E. Gesser (eds.), NASA/Astronomical
Data Center, Goddard Space Flight Center, Greenbelt, MD
Graham, J. A. 1982, PASP, 94, 244
Grevesse, N., & Sauval, A. J. 1998, Space Science Reviews, 85, 161
Hambly, N. C., Smartt, S. J., Hodgkin, S. T., Jameson, R. F., Kemp, S. N., Rolleston,
W. R. J., & Steele, I. A. 1999, MNRAS, 309, L33
Henry, T. J., Walkowicz, L. M., Barto, T. C., & Golimowski, D. A. 2002, AJ, 123, 2002
Henry, T. J., Backman, D. E., Blackwell, J., Okimura, T., & Jue, S. 2003, The Future of
Small Telescopes In The New Millennium. Volume III - Science in the Shadow of
Giants, 111
Henry, T. J., Subasavage, J. P., Brown, M. A., Beaulieu, T. D., Jao, W., & Hambly, N. C.
2004, AJ, 128, 2460
Henry, T. J., Jao, W.-C., Subasavage, J. P., Beaulieu, T. D., Ianna, P. A., Costa, E., &
Méndez, R. A. 2006, AJ, 132, 2360
Holberg, J. B., Oswalt, T. D., & Sion, E. M. 2002, ApJ, 571, 512
– 15 –
Holberg, J. B., & Bergeron, P. 2006, AJ, 132, 1223
Iben, I., & Renzini, A. 1984, Phys. Rep., 105, 329
Jao, W.-C., Henry, T. J., Subasavage, J. P., Brown, M. A., Ianna, P. A., Bartlett, J. L.,
Costa, E., & Méndez, R. A. 2005, AJ, 129, 1954
Kawka, A., & Vennes, S. 2006, ApJ, 643, 402
Kilic, M., von Hippel, T., Leggett, S. K., & Winget, D. E. 2006, ApJ, 646, 474
Kilkenny, D., O’Donoghue, D., Koen, C., Stobie, R. S., & Chen, A. 1997, MNRAS, 287, 867
Kleinman, S. J., et al. 2004, ApJ, 607, 426
Landolt, A. U. 1992, AJ, 104, 340
Lépine, S., Shara, M. M., & Rich, R. M. 2003, AJ, 126, 921
Lépine, S., & Shara, M. M. 2005, AJ, 129, 1483
Lépine, S., Rich, R. M., & Shara, M. M. 2005, ApJ, 633, L121
Liebert, J., Bergeron, P., & Holberg, J. B. 2003, AJ, 125, 348
Liebert, J., Bergeron, P., & Holberg, J. B. 2005, ApJS, 156, 47
Luyten, W. J. 1949, ApJ, 109, 528
Luyten, W. J. 1979, LHS Catalogue (2nd ed.; Minneapolis: Univ. of Minnesota Press)
Luyten, W. J. 1979, New Luyten Catalogue of Stars with Proper Motions Larger than Two
Tenths of an Arcsecond (Minneapolis: Univ. of Minnesota Press)
McCook, G. P., & Sion, E. M. 1999, ApJS, 121, 1
Oppenheimer, B. R., Hambly, N. C., Digby, A. P., Hodgkin, S. T., & Saumon, D. 2001,
Science, 292, 698
Press, W. H., Teukolsky, S. A., Vetterling, W. T., & Flannery, B. P. 1992, Numerical Recipes
in FORTRAN, 2nd edition (Cambridge: Cambridge University Press), 644
Provencal, J. L., Shipman, H. L., Hog, E., & Thejll, P. 1998, ApJ, 494, 759
Pokorny, R. S., Jones, H. R. A., Hambly, N. C., & Pinfield, D. J. 2004, A&A, 421, 763
– 16 –
Salim, S., Rich, R. M., Hansen, B. M., Koopmans, L. V. E., Oppenheimer, B. R., & Bland-
ford, R. D. 2004, ApJ, 601, 1075
Scholz, R.-D., Szokoly, G. P., Andersen, M., Ibata, R., & Irwin, M. J. 2002, ApJ, 565, 539
Skrutskie, M. F., et al. 2006, AJ, 131, 1163
Smart, R. L., et al. 2003, A&A, 404, 317
Subasavage, J. P., Henry, T. J., Hambly, N. C., Brown, M. A., & Jao, W. 2005, AJ, 129, 413
Subasavage, J. P., Henry, T. J., Hambly, N. C., Brown, M. A., Jao, W.-C., & Finch, C. T.
2005, AJ, 130, 1658
This preprint was prepared with the AAS LATEX macros v5.2.
Table 1. Optical and Infrared Photometry, and Derived Parameters for New and Known White Dwarfs.
WD VJ RC IC # J σJ H σH KS σK
Teff Comp Dist SpT Notes
Name Obs (K) (pc)
New Spectroscopically Confirmed White Dwarfs
0034−602............. 14.08 14.19 14.20 3 14.37 0.04 14.55 0.06 14.52 0.09 14655±1413 H 35.8±5.7 DA3.5
0121−429............. 14.83 14.52 14.19 4 13.85 0.02 13.63 0.04 13.53 0.04 6369± 137 H · · · ± · · · DAH a
0216−398............. 15.75 15.55 15.29 3 15.09 0.04 14.83 0.06 14.89 0.14 7364± 241 H 29.9±4.7 DA7.0
0253−755............. 16.70 16.39 16.08 2 15.77 0.07 15.76 0.15 15.34 null 6235± 253 He 34.7±5.5 DC
0310−624............. 15.92 15.99 16.03 2 16.13 0.10 16.31 0.27 16.50 null 13906±1876 H · · · ± · · · DA3.5 b
0344+014............. 16.52 16.00 15.54 2 15.00 0.04 14.87 0.09 14.70 0.12 5084± 91 He 19.9±3.1 DC
0404−510............. 15.81 15.76 15.70 2 15.74 0.06 15.55 0.13 15.59 null 10052± 461 H 53.5±8.5 DA5.0
0501−555............. 16.35 16.17 15.98 2 15.91 0.08 15.72 0.15 15.82 0.26 7851± 452 He 44.8±6.9 DC
0511−415............. 16.00 15.99 15.93 2 15.96 0.08 15.97 0.15 15.20 null 10393± 560 H 61.8±10.8 DA5.0
0525−311............. 15.94 16.03 16.03 2 16.20 0.12 16.21 0.25 14.98 null 12941±1505 H 76.3±13.6 DA4.0
0607−530............. 15.99 15.92 15.78 3 15.82 0.07 15.66 0.14 15.56 0.21 9395± 426 H 51.7±9.0 DA5.5
0622−329............. 15.47 15.41 15.36 2 15.44 0.06 15.35 0.11 15.53 0.25 · · · ± · · · · · · · · · ± · · · DAB c
0821−669............. 15.34 14.82 14.32 3 13.79 0.03 13.57 0.03 13.34 0.04 5160± 95 H 11.5±1.9 DA10.0
0840−136............. 15.72 15.36 15.02 3 14.62 0.03 14.42 0.05 14.54 0.09 · · · ± · · · · · · · · · ± · · · DZ d
1016−308............. 14.67 14.75 14.81 2 15.05 0.04 15.12 0.08 15.41 0.21 16167±1598 H 50.6±9.2 DA3.0
1054−226............. 16.02 15.82 15.62 2 15.52 0.05 15.40 0.11 15.94 0.26 8266± 324 H 41.0±7.0 DA6.0 e
1105−340............. 13.66 13.72 13.79 2 13.95 0.03 13.98 0.04 14.05 0.07 13926± 988 H 28.2±4.8 DA3.5 f
1149−272............. 15.87 15.59 15.37 4 15.17 0.05 14.92 0.06 14.77 0.11 6188± 194 He (+C) 24.0±3.8 DQ
1243−123............. 15.57 15.61 15.64 2 15.74 0.07 15.73 0.11 16.13 null 12608±1267 H 62.6±10.7 DA4.0
1316−215............. 16.67 16.33 15.99 2 15.56 0.05 15.33 0.08 15.09 0.14 6083± 201 H 31.6±5.3 DA8.5
1436−781............. 16.11 15.82 15.49 2 15.04 0.04 14.88 0.08 14.76 0.14 6246± 200 H 26.0±4.3 DA8.0
1452−310............. 15.85 15.77 15.63 2 15.58 0.06 15.54 0.09 15.50 0.22 9206± 375 H 46.8±8.1 DA5.5
1647−327............. 16.21 15.85 15.49 3 15.15 0.05 14.82 0.08 14.76 0.11 6092± 193 H 25.5±4.2 DA8.5
1742−722............. 15.53 15.62 15.70 2 15.85 0.08 15.99 0.18 15.65 null 15102±2451 H 71.7±12.9 DA3.5
1946−273............. 14.19 14.31 14.47 2 14.72 0.04 14.77 0.09 14.90 0.13 21788±3304 H 52.0±9.9 DA2.5
2008−600............. 15.84 15.40 14.99 4 14.93 0.05 15.23 0.11 15.41 null 5078± 221 He · · · ± · · · DC g
2008−799............. 16.35 15.96 15.57 3 15.11 0.04 15.03 0.08 14.64 0.09 5807± 161 H 24.5±4.1 DA8.5
2035−369............. 14.94 14.85 14.72 2 14.75 0.04 14.72 0.06 14.84 0.09 9640± 298 H 33.1±5.7 DA5.0
2103−397............. 15.31 15.15 14.91 2 14.79 0.03 14.63 0.04 14.64 0.08 7986± 210 H 28.2±4.8 DA6.5
2138−332............. 14.47 14.30 14.16 3 14.17 0.03 14.08 0.04 13.95 0.06 7188± 291 He (+Ca) 17.3±2.7 DZ
2157−574............. 15.96 15.73 15.49 3 15.18 0.04 15.05 0.07 15.28 0.17 7220± 246 H 32.0±5.4 DAZ
2218−416............. 15.36 15.35 15.24 2 15.38 0.04 15.14 0.09 15.39 0.15 10357± 414 H 45.6±8.0 DA5.0
2231−387............. 16.02 15.88 15.62 2 15.57 0.06 15.51 0.11 15.11 0.15 8155± 336 H 40.6±6.9 DA6.0
Known White Dwarfs without a Trigonometric Parallax Estimated to be Within 25 pc
0141−675 ............ 13.82 13.52 13.23 3 12.87 0.02 12.66 0.03 12.58 0.03 6484± 128 H 9.7±1.6 DA8.0
0806−661 ............ 13.73 13.66 13.61 3 13.70 0.02 13.74 0.03 13.78 0.04 10753± 406 He 21.1±3.5 DQ
1009−184 ............ 15.44 15.18 14.91 3 14.68 0.04 14.52 0.05 14.31 0.07 6449± 194 He 20.9±3.2 DZ h
1036−204 ............ 16.24 15.54 15.34 3 14.63 0.03 14.35 0.04 14.03 0.07 4948± 70 He 16.2±2.5 DQ i
1202−232 ............ 12.80 12.66 12.52 3 12.40 0.02 12.30 0.03 12.34 0.03 8623± 168 H 10.2±1.7 DA6.0
1315−781 ............ 16.16 15.73 15.35 2 14.89 0.04 14.67 0.08 14.58 0.12 5720± 162 H 21.6±3.6 DC j
1339−340 ............ 16.43 16.00 15.56 2 15.00 0.04 14.75 0.06 14.65 0.10 5361± 138 H 21.2±3.5 DA9.5
1756+143 ............ 16.30 16.12 15.69 1 14.93 0.04 14.66 0.06 14.66 0.08 5466± 151 H 22.4±3.4 DA9.0 k
Table 1—Continued
WD VJ RC IC # J σJ H σH KS σK
Teff Comp Dist SpT Notes
Name Obs (K) (pc)
1814+134 ............ 15.85 15.34 14.86 2 14.38 0.04 14.10 0.06 14.07 0.06 5313± 115 H 15.6±2.5 DA9.5
2040−392 ............ 13.74 13.77 13.68 2 13.77 0.02 13.82 0.03 13.81 0.05 10811± 325 H 23.1±4.0 DA4.5
2211−392 ............ 15.91 15.61 15.24 2 14.89 0.03 14.64 0.05 14.56 0.08 6243± 167 H 23.5±4.0 DA8.0
2226−754A........... 16.57 15.93 15.33 2 14.66 0.04 14.66 0.06 14.44 0.08 4230± 104 H 12.8±2.0 DC l
2226−754B........... 16.88 16.17 15.51 2 14.86 0.04 14.82 0.06 14.72 0.12 4177± 112 H 14.0±2.2 DC l
Known White Dwarfs without a Trigonometric Parallax Estimated to be Beyond 25 pc
0024−556............. 15.17 15.15 15.07 2 15.01 0.04 15.23 0.10 15.09 0.14 10007± 378 H 39.8±6.8 DA5.0
0150+256............. 15.70 15.52 15.33 2 15.07 0.04 15.07 0.09 15.15 0.14 7880± 280 H 33.0±5.6 DA6.5
0255−705 ............ 14.08 14.03 14.00 2 14.04 0.03 14.12 0.04 13.99 0.06 10541± 326 H 25.8±4.5 DA5.0
0442−304............. 16.03 15.93 15.86 2 15.94 0.09 15.81 null 15.21 null 9949± 782 He 55.1±9.1 DQ
0928−713 ............ 15.11 14.97 14.83 3 14.77 0.03 14.69 0.06 14.68 0.09 8836± 255 H 30.7±5.3 DA5.5
1143−013............. 16.39 16.08 15.79 1 15.54 0.06 15.38 0.08 15.18 0.16 6824± 250 H 34.4±5.8 DA7.5
1237−230 ............ 16.53 16.13 15.74 2 15.35 0.05 15.08 0.08 14.94 0.11 5841± 173 H 26.9±4.5 DA8.5
1314−153............. 14.82 14.89 14.97 2 15.17 0.05 15.26 0.09 15.32 0.21 15604±2225 H 52.7±9.5 DA3.0
1418−088 ............ 15.39 15.21 15.01 2 14.76 0.04 14.73 0.06 14.76 0.10 7872± 243 H 28.5±4.8 DA6.5
1447−190............. 15.80 15.59 15.32 2 15.06 0.04 14.87 0.07 14.78 0.11 7153± 235 H 29.1±4.9 DA7.0
1607−250............. 15.19 15.12 15.09 2 15.08 0.08 15.08 0.08 15.22 0.15 10241± 457 H 41.2±7.2 DA5.0
aDistance via SED fit (not listed) is underestimated because object is likely an unresolved double degenerate with one magnetic component (see § 4.2). Instead, we
adopt the trigonometric parallax distance of 17.7 ± 0.7 pc derived via CTIOPI.
bDistance via SED fit (not listed) is underestimated because object is likely a distant (well beyond 25 pc) unresolved double degenerate (see § 4.2).
cDistance via SED fit (not listed) is underestimated because object is likely a distant (well beyond 25 pc) unresolved double degenerate with components of type DA
and DB (see § 4.2). Temperatures derived from the spectroscopic fit yield 9,640 ± 303 K and 14,170 ± 1,228 K for the DA and DB respectively.
dObject is likely cooler than Teff ∼5000 K and the theoretical models do not provide an accurate treatment at these temperatures (see § 4.2). Instead, we use the
linear photometric distance relation of Salim et al. (2004) and obtain a distance estimate of 19.3 ± 3.9 pc.
eThis object was observed as part of the Edinburgh-Cape survey and was classified as a sdB+ (Kilkenny et al. 1997).
fDistance of 19.1 ± 3.0 pc is estimated using V RIJHKS for the common proper motion companion M dwarf and the relations of Henry et al. (2004). System is
possibly within 25 pc. (see § 4.2).
gDistance estimate is undetermined. Instead, we adopt the distance measured via trogonometric parallax of 17.1 ± 0.4 pc (see § 4.2).
hNot listed in McCook & Sion (1999) but identified as a DC/DQ WD by Henry et al. (2002). We obtained blue spectra that show Ca II H & K absorption and classify
this object as a DZ.
iThe SED fit to the photometry is marginal. This object displays deep swan band absorption that significantly affects its measured magnitudes.
jNot listed in McCook & Sion (1999) but identified as a WD by Luyten (1949). Spectral type is derived from our spectra.
kAs of mid-2004, object has moved onto a background source. Photometry is probably contaminated, which is consistent with the poor SED fit for this object.
lSpectral type was determined using spectra published by Scholz et al. (2002).
– 19 –
Table 2. Distance Estimate Statistics for New and Known White Dwarfs.
Proper motion d ≤ 10 pc 10 pc < d ≤ 25 pc d > 25 pc
µ ≥ 1.′′0 yr−1......................... 1 6 1
1.′′0 yr−1 > µ ≥ 0.′′8 yr−1...... 0 0 0
0.′′8 yr−1 > µ ≥ 0.′′6 yr−1...... 0 2 2
0.′′6 yr−1 > µ ≥ 0.′′4 yr−1...... 0 6 11
0.′′4 yr−1 > µ ≥ 0.′′18 yr−1.... 0 5 22
Total.................................... 1 19 36
– 20 –
Table 3. Astrometry and Alternate Designations for New and Known White Dwarfs.
WD Name RA Dec PM PA Ref Alternate Names
(J2000.0) (J2000.0) (arcsec yr−1) (deg)
New Spectroscopically Confirmed White Dwarfs
0034−602......... 00 36 22.31 −59 55 27.5 0.280 069.0 L NLTT 1993 = LP 122-4 = · · ·
0121−429......... 01 24 03.98 −42 40 38.5 0.538 155.2 L LHS 1243 = NLTT 4684 = LP 991-16
0216−398......... 02 18 31.51 −39 36 33.2 0.500 078.6 L LHS 1385 = NLTT 7640 = LP 992-99
0253−755......... 02 52 45.64 −75 22 44.5 0.496 063.5 S SCR 0252-7522 = · · · = · · ·
0310−624......... 03 11 21.34 −62 15 15.7 0.416 083.3 S SCR 0311-6215 = · · · = · · ·
0344+014......... 03 47 06.82 +01 38 47.5 0.473 150.4 S LHS 5084 = NLTT 11839 = LP 593-56
0404−510......... 04 05 32.86 −50 55 57.8 0.320 090.7 P LEHPM 1-3634 = · · · = · · ·
0501−555......... 05 02 43.43 −55 26 35.2 0.280 191.9 P LEHPM 1-3865 = · · · = · · ·
0511−415......... 05 13 27.80 −41 27 51.7 0.292 004.4 P LEHPM 2-1180 = · · · = · · ·
0525−311......... 05 27 24.33 −31 06 55.7 0.379 200.7 P NLTT 15117 = LP 892-45 = LEHPM 2-521
0607−530......... 06 08 43.81 −53 01 34.1 0.246 327.6 P LEHPM 2-2008 = · · · = · · ·
0622−329......... 06 24 25.78 −32 57 27.4 0.187 177.7 P LEHPM 2-5035 = · · · = · · ·
0821−669......... 08 21 26.70 −67 03 20.1 0.758 327.6 S SCR 0821-6703 = · · · = · · ·
0840−136......... 08 42 48.45 −13 47 13.1 0.272 263.0 S NLTT 20107 = LP 726-1 = · · ·
1016−308......... 10 18 39.84 −31 08 02.0 0.212 304.0 L NLTT 23992 = LP 904-3 = LEHPM 2-5779
1054−226......... 10 56 38.64 −22 52 55.9 0.277 349.7 P NLTT 25792 = LP 849-31 = LEHPM 2-1372
1105−340......... 11 07 47.89 −34 20 51.4 0.287 168.0 S SCR 1107-3420A = · · · = · · ·
1149−272......... 11 51 36.10 −27 32 21.0 0.199 278.3 P LEHPM 2-4051 = · · · = · · ·
1243−123......... 12 46 00.69 −12 36 19.9 0.406 305.4 S SCR 1246-1236 = · · · = · · ·
1316−215......... 13 19 24.72 −21 47 55.0 0.467 179.2 S NLTT 33669 = LP 854-50 = WT 2034
1436−781......... 14 42 51.54 −78 23 53.6 0.409 272.0 S NLTT 38003 = LP 40-109 = LTT 5814
1452−310......... 14 55 23.47 −31 17 06.4 0.199 174.2 P LEHPM 2-4029 = · · · = · · ·
1647−327......... 16 50 44.32 −32 49 23.2 0.526 193.8 L LHS 3245 = NLTT 43628 = LP 919-1
1742−722......... 17 48 31.21 −72 17 18.5 0.294 228.2 P LEHPM 2-1166 = · · · = · · ·
1946−273......... 19 49 19.78 −27 12 25.7 0.213 162.0 L NLTT 48270 = LP 925-53 = · · ·
2008−600......... 20 12 31.75 −59 56 51.5 1.440 165.6 S SCR 2012-5956 = · · · = · · ·
2008−799......... 20 16 49.66 −79 45 53.0 0.434 128.4 S SCR 2016-7945 = · · · = · · ·
2035−369......... 20 38 41.42 −36 49 13.5 0.230 104.0 L NLTT 49589 = L 495-42 = LEHPM 2-3290
2103−397......... 21 06 32.01 −39 35 56.7 0.266 151.7 P LEHPM 2-1571 = · · · = · · ·
2138−332......... 21 41 57.56 −33 00 29.8 0.210 228.5 P NLTT 51844 = L 570-26 = LEHPM 2-3327
2157−574......... 22 00 45.37 −57 11 23.4 0.233 252.0 P LEHPM 1-4327 = · · · = · · ·
2218−416......... 22 21 25.37 −41 25 27.0 0.210 143.4 P LEHPM 1-4598 = · · · = · · ·
2231−387......... 22 33 54.47 −38 32 36.9 0.370 220.5 P NLTT 54169 = LP 1033-28 = LEHPM 1-4859
Known White Dwarfs without a Trigonometric Parallax Estimated to be Within 25 pc
0141−675 ........ 01 43 00.98 −67 18 30.3 1.048 197.8 L LHS 145 = NLTT 5777 = L 88-59
0806−661 ........ 08 06 53.76 −66 18 16.6 0.454 131.4 S NLTT 19008 = L 97-3 = · · ·
1009−184 ........ 10 12 01.88 −18 43 33.2 0.519 268.2 S WT 1759 = LEHPM 2-220 = · · ·
1036−204 ........ 10 38 55.57 −20 40 56.7 0.628 330.3 L LHS 2293 = NLTT 24944 = LP 790-29
1202−232 ........ 12 05 26.66 −23 33 12.1 0.227 002.0 L NLTT 29555 = LP 852-7 = LEHPM 2-1894
1315−781 ........ 13 19 25.63 −78 23 28.3 0.477 139.2 S NLTT 33551 = L 40-116 = · · ·
1339−340 ........ 13 42 02.88 −34 15 19.4 2.547 296.7 Le PM J13420-3415 = · · · = · · ·
1756+143 ........ 17 58 22.90 +14 17 37.8 1.014 235.4 Le LSR 1758+1417 = · · · = · · ·
1814+134 ........ 18 17 06.48 +13 28 25.0 1.207 201.5 Le LSR 1817+1328 = · · · = · · ·
2040−392 ........ 20 43 49.21 −39 03 18.0 0.306 179.0 L NLTT 49752 = L 495-82 = · · ·
2211−392 ........ 22 14 34.75 −38 59 07.3 1.056 110.1 O WD J2214-390 = LEHPM 1-4466 = · · ·
2226−754A........ 22 30 40.00 −75 13 55.3 1.868 167.5 S SSSPM J2231-7514 = · · · = · · ·
2226−754B........ 22 30 33.55 −75 15 24.2 1.868 167.5 S SSSPM J2231-7515 = · · · = · · ·
Known White Dwarfs without a Trigonometric Parallax Estimated to be Beyond 25 pc
0024−556......... 00 26 40.69 −55 24 44.1 0.580 211.8 L LHS 1076 = NLTT 1415 = L 170-27
0150+256......... 01 52 51.93 +25 53 40.7 0.220 076.0 L NLTT 6275 = G 94-21 = · · ·
0255−705......... 02 56 17.22 −70 22 10.8 0.682 097.9 L LHS 1474 = NLTT 9485 = L 54-5
0442−304......... 04 44 29.38 −30 21 14.2 0.196 199.5 P NLTT 13882 = LP 891-65 = HE 0442-3027
0928−713......... 09 29 07.97 −71 33 58.8 0.439 320.2 S NLTT 21957 = L 64-40 = · · ·
1143−013......... 11 46 25.77 −01 36 36.8 0.563 140.2 S LHS 2455 = NLTT 28493 = · · ·
– 21 –
Table 3—Continued
WD Name RA Dec PM PA Ref Alternate Names
(J2000.0) (J2000.0) (arcsec yr−1) (deg)
1237−230......... 12 40 24.18 −23 17 43.8 1.102 219.9 L LHS 339 = NLTT 31473 = LP 853-15
1314−153......... 13 16 43.59 −15 35 58.3 0.708 196.7 L LHS 2712 = NLTT 33503 = LP 737-47
1418−088......... 14 20 54.93 −09 05 08.7 0.480 266.8 S LHS 5270 = NLTT 37026 = · · ·
1447−190......... 14 50 11.93 −19 14 08.7 0.253 285.4 P NLTT 38499 = LP 801-14 = LEHPM 2-1835
1607−250......... 16 10 50.21 −25 13 16.0 0.209 314.0 L NLTT 42153 = LP 861-31 = · · ·
References. — (L) Luyten 1979a,b, (Le) Lépine et al. 2003, Lépine et al. 2005, (O) Oppenheimer et al. 2001, (P) Pokorny et al.
2004, (S) Subasavage et al. 2005a,b, this work
– 22 –
Fig. 1.— Reduced proper motion diagram used to select WD candidates for spectroscopic
follow-up. Plotted are the new high proper motion objects from Subasavage et al. (2005a,b).
The line is a somewhat arbitrary boundary between the WDs (below) and the subdwarfs
(just above). Main sequence dwarfs fall above and to the right of the subdwarfs, although
there is significant overlap. Asterisks indicate the 33 new WDs reported here. Three dots
in the WD region are deferred to a future paper. The point labeled “sd” is a confirmed
subdwarf contaminant of the WD sample.
Fig. 2.— Spectral plots of the hot (Teff ≥ 10000 K) DA WDs from the new sample, plotted
in descending Teff as derived from the SED fits to the photometry. Note that some of the
flux calibrations are not perfect, in particular, at the blue end.
Fig. 3.— Spectral plots of cool (Teff < 10000 K) DA WDs from the new sample, plotted in
descending Teff as derived from the SED fits to the photometry. Note that some of the flux
calibrations are not perfect, in particular, at the blue end.
Fig. 4.— Spectral plots of the four featureless DC white dwarfs from the new sample, plotted
in descending Teff as derived from the SED fits to the photometry. Note that some of the
flux calibrations are not perfect, in particular, at the blue end.
Fig. 5.— Comparison plot of the values of Teff derived from photometric SED fitting vs
those derived from spectral fitting for 25 of the DA WDs in the new sample. The solid line
represents equal temperatures. The elevated point, 0310−624, is discussed in § 4.2.
Fig. 6.— Spectral plot of WD 0121−429. The inset plot displays the spectrum (light line)
in the Hα region to which a magnetic fit (heavy line), as outlined in Bergeron et al. (1992a),
was performed using the Teff obtained from the SED fit to the photometry. The resulting
magnetic parameters are listed below the fit.
Fig. 7.— Spectral plot of WD 0622−329. The inset plot displays the spectrum (light line)
in the region to which the model (heavy line) was fit assuming the spectrum is a convolution
of a DB component and a slightly cooler DA component. Best fit physical parameters are
listed below the fit for each component.
Fig. 8.— (top panel) Spectral plot of WD 0840−136. The DZ model failed to reproduce
the spectrum presumably because this object is cooler than Teff ∼ 5000 K where additional
pressure effects, not included in the model, become important. (bottom panel) Spectral plot
of WD 2138−332. The inset plot displays the spectrum (light line) in the region to which
the model (heavy line) was fit.
– 23 –
Fig. 9.— Spectral plot of WD 1149−272. The inset plot displays the spectrum (light line)
in the region to which the model (heavy line) was fit.
Fig. 10.— Spectral energy distribution plot of WD 2008−600 with the distance constrained
by the trigonometric distance of 17.1 ± 0.4 pc. Best fit physical parameters are listed below
the fit. Points are fit values; error bars are derived from the uncertainties in the magnitudes
and the parallax.
– 24 –
– 25 –
– 26 –
– 27 –
– 28 –
– 29 –
– 30 –
– 31 –
– 32 –
– 33 –
	Introduction
	Candidate Selection
	Data and Observations
	Astrometry and Nomenclature
	Spectroscopy
	Photometry
	Analysis
	Modeling of Physical Parameters
	Comments on Individual Systems
	Discussion
	Acknowledgments
	Appendix
