Issue |
A&A
Volume 508, Number 1, December II 2009
|
|
---|---|---|
Page(s) | 339 - 344 | |
Section | Stellar structure and evolution | |
DOI | https://doi.org/10.1051/0004-6361/200912094 | |
Published online | 01 October 2009 |
A&A 508, 339-344 (2009)
A hot white dwarf luminosity function from the Sloan Digital Sky Survey
J. Krzesinski1,2 - S. J. Kleinman3 - A. Nitta3 - S. Hügelmeyer4 - S. Dreizler4 - J. Liebert5 - H. Harris6
1 - Mt. Suhora Observatory, Cracow Pedagogical University, ul. Podchorazych 2, 30-084 Cracow, Poland
2 -
Apache Point Observatory, PO Box 59, Sunspot, NM 88349, USA
3 -
Gemini Observatory, 670 N. A'Ohoku Place, Hilo, HI 96720, USA
4 -
Institut für Astrophysik, Georg-August-Universität Göttingen, Friedrich-Hund-Platz 1, 37077 Göttingen, Germany
5 -
Steward Observatory, University of Arizona, Tucson, AZ 85726, USA
6 -
US Naval Observatory, PO Box 1149, Flagstaff, AZ 86002-1149, USA
Received 17 March 2009 / Accepted 9 September 2009
Abstract
Aims. We present a hot white dwarf (WD) luminosity function
(LF) using data taken from the Sloan Digital Sky Survey (SDSS) Data
Release 4. We present and discuss a combined LF, along with separate DA
and non-DA as LFs. We explore the completeness of our LFs and interpret
a sudden drop in the non-DA LF near 2
as a transition of the non-DA WD atmosphere into the DA one during WD evolution. Our LF extends roughly between
or equivalently,
K.
Our LF should now be useful for estimates of recent star formation and
for studies of neutrino and other potential particle emission losses in
hot WDs.
Methods. To create a sample whose completeness can be
characterized fully, we used stars whose spectra were obtained via
the SDSS's ``hot standard'' target selection criteria. The hot standard
stars were purposefully targeted to a high level of completeness by the
SDSS for calibration purposes. We are fortunate that many of them are
hot white dwarfs stars. We further limited the sample to stars with
fitted temperatures exceeding 23 500 K and
.
We determined stellar distances for our sample based on their absolute
SDSS g filter magnitudes, derived from WD stellar atmosphere model fits
to the SDSS stellar spectra.
Results. We compared our LF with those of other researchers
where overlap occurs; however, our LFs are unique in their extension to
the most luminous/hottest WDs. The cool end of our LF connects with the
hot end of previously determined SDSS WD LFs and agreement here is
quite good. It is also good with previous non-SDSS WD LFs. We note
distinct differences between the DA and non-DA LFs and discuss the
reliability of the DA LF at its hot end. We have extended the range of
luminosities covered in the most recent WD LFs. The SDSS sample is
understood quite well and its exploration should contribute to a number
of new insights into early white dwarf evolution.
Key words: stars: luminosity function, mass function - white dwarfs
1 Introduction
The luminosity function (LF) and space density of WDs provide constraints on the local star formation rate and history of the Galactic disk in the Solar neighborhood. They can also be used to help determine the role of weakly interacting particles in WD cooling (Irsen et al. 2008) and thus their ages. A large WD sample of known completeness is required to determine a precise LF to address these questions. Luckily, the Sloan Digital Sky Survey (SDSS) (Pier 2003; Gunn et al. 2006; Stoughton et al. 2002; Gunn et al. 1998; Tucker et al. 2006; York et al. 2000) provides the perfect sample for this effort. Although primarily designed as an extra-Galactic survey, the SDSS is an extremely rich source of stellar (and in particular, white dwarf) spectra. Eisenstein et al. (2006a) published a catalog of spectroscopically confirmed white dwarf stars from the SDSS Data Release 4 (DR4: Adelman-McCarthy et al. 2006). Their catalog contained 9316 WDs, nearly 3.6 times the number of WDs found in the first SDSS WD catalog (Kleinman et al. 2004) based on DR1 (Abazajian et al. 2003) data, consistent with the ratio of spectroscopic sky coverage (4783 / 1360 square degrees) in DR4 vs. DR1.
Harris et al. (2006) have constructed a reduced proper-motion based
WD LF using the photometric sample from SDSS DR3 (Abazajian et al. 2005)
and proper motions from
Monet et al. (2003). The large data set available from SDSS enabled them to
sample the LF within a magnitude range of
using 0.5 mag bins, with each bin containing hundreds of stars for quite
excellent statistics.
Hu et al. (2007) also calculated a DA WD LF based on the much smaller SDSS DR1
WD sample, but extending to
at its blue end.
These LFs compare well with that of Liebert et al. (2005), which is based on DAs
from the Palomar Green survey (Green et al. 1986). Most recently,
DeGennaro et al. (2008) have presented a DA WD LF also based on SDSS DR4 data,
using
statistical arguments and a comparison to the Harris et al. (2006) LF to estimate
sample completeness. The agreements and disagreements between the two LFs
are fairly well understood, although their cutoff at
offers only a small chance of comparision to our own LF,
which extends to much hotter stars.
Prior to the flurry of new SDSS-based WD LFs, WD LFs were constructed by
for example Flemming et al. (1986) and Vennes et al. (2002). There were also
attempts to estimate the local non-DA WDs space density dating back
to 1985 (Wesemael et al. 1985). None, however, extend below
mag, while our new LF extends to below
mag.
Traditionally, distances in WD LFs have been calculated
using a statistical proper motion method
(e.g. Oswalt et al. 1996; Liebert et al. 1988; Evans 1992, and others).
For our relatively bright, therefore distant stars, this approach
is less likely to work, since many are simply too far away to have
measurable proper motions.
We therefore calculated bolometric magnitudes derived from fitting
theoretical atmosphere model to the SDSS spectra and compared them
to the observed SDSS g filter
(SDSS filters: Fukugita et al. 1996)
magnitudes to obtain distances to each of our objects.
Unlike all the other WDLFs, ours is based solely on
spectroscopically-confirmed WDs from the SDSS spectroscopic
sample and covers a
range from -0.75-7.0.
2 Our hot white dwarf sample
The SDSS ``hot_std'' category targets bright (g<19), isolated, hot stars for
calibration purposes. In short, the following criteria define the
hot_stds, where the subscript
indicates a fully-dereddened
magnitude:
- g > 14
-
-
-
.
![]() |
Figure 1:
The photometric
|
Open with DEXTER |
Figure 1 is a plot of hot DAs and DBs from Eisenstein et al. (2006a)
along with the hot_std
and
color cuts.
The
color selection ends up cutting out DAs at cooler temperatures than the DBs.
Objects above the horizontal line (some DAs cooler than 23 500 K) are
not targeted as hot_std stars, thereby introducing some incompleteness
in our sample below that temperature.
We therefore imposed an additional cut on our sample of
(the vertical line in
Fig. 1), so the temperature
selection of DAs and non-DAs in our sample is unaffected by the DAs higher
(u-g) values.
3 Luminosity function
Once our white dwarf sample was well-defined, we determined their distances and absolute bolometric magnitudes. Next, we calculated the sample's completeness and volume corrections before forming our final luminosity functions.
3.1 The atmospheric models and distances
We used the Eisenstein et al. (2006a) fit
and
parameters for the pure DA and DB stars in
our sample.
Unfortunately, Eisenstein et al. (2006a) do not include any predicted
absolute photometric magnitudes in their output, so we needed
additional work to convert the
and
fits into an absolute
magnitude which we could use to calculate distances.
Bergeron (2001, private communication) calculated a grid of
,
vs.
and absolute SDSS u, g, r, i, z filters
magnitudes for pure DA and DB WDs.
Eisenstein et al. (2006a) demonstrated that their models (from Detlev Koester; published
in Finley et al. 1997) are in good agreement with Bergeron's, so we can
use the
and
fits from their Koester model fits to
obtain absolute SDSS filter magnitudes from the Bergeron grid. We
used these absolute magnitudes to compare with the observed magnitudes
and obtain distances to all our objects.
For the non DA/DB stars, we used other spectral models and fits accordingly: Rauch & Deetjen (2003) non-LTE H+He models for the DAOs and Hügelmeyer et al. (2006) models for DO and PG1159 spectra. We used stellar radii taken from evolutionary models of Wood (1995) with thick and thin hydrogen layers for DAO and DO WDs and Bloecker (1995), Schoenberner (1983) and Wood & Faulkner (1986) models for PG1159 stars. DAB (3 stars) and DBA (1 star) WDs were treated in the same way as DA and DB WDs, respectively. For magnetic or binary DAs and DBs the fit parameters were taken from the Eisenstein et al. (2006a) DR4 WD catalog; for binary DOs, from Hügelmeyer et al. (2006). Finally we calculated two LFs, one for all WDs including magnetic and binary WDs, and one without them. The LFs with and without the magnetic and binary WDs are not significantly different from each other, so we therefore continue with these objects included in our further analysis.
In general, an absolute SDSS filter magnitude
can be obtained by
implementing the modified Girardi et al. (2002) formula (Eq. (2)):

where:
flt | one of u, g, r, i, z SDSS filters, |
![]() | AB system zero magnitude for |
a SDSS filter | |
Oke & Gunn (1983), | |
c | speed of light in cm/s, |
A2 cm | Angstrom to cm conversion, |
![]() | 1 parsec, |
![]() | solar radius in cm, |
![]() | flux density from the model, |
Filter![]() | SDSS filter transmission within |
the wavelength range
![]() ![]() |
|
![]() | wavelength in Angstroms, |
![]() | stellar radius taken from |
evolutionary models | |
zeropoint(flt) [mag] | SDSS photometric zero point. |
Here zeropoint(u) = -0.04, zeropoint(g) = zeropoint(r) = 0.00, zeropoint(i) = 0.015 and zeropoint(z) = 0.03. The SDSS magnitude system (photometric system: Hogg 2001; Ivezic et al. 2004; Smith et al. 2002) is close to being on the AB system Oke & Gunn (1983), but not quite, so the above zero points are needed (Abazajian et al. 2004; Eisenstein et al. 2006b).
We used the Girardi et al. (2002) equations, above, to obtain absolute SDSS g filter
magnitudes for all WDs in our sample. For most of them, we calculated
absolute bolometric magnitudes (by extending the wavelength range
in the equation from 0 to infinity). For some of the DOs and PG1159 stars
in the sample, however, we do not have complete enough models to
estimate a total stellar flux, so in those cases we used the standard
equation, below, to obtain absolute bolometric magnitudes:

We compared results using this equation with the absolute magnitudes determined by our model grids and found the answers to be in good agreement (

We do not have models for the hottest DBs in our sample those between about 30 000 K and 42 500 K, the so-called ``DB gap'' region (Eisenstein et al. 2006a,b). For these objects, therefore, we extended the Bergeron grid using the DO models from Hügelmeyer et al. (2006) as a guide. However, this procedure likely introduces magnitude errors of approximately 0.5 mag.
Finally, we had to apply a reddening correction to our calculated absolute and observed magnitudes in order to obtain distances to our objects. Since we don't know how much of the measured interstellar reddening is really between us and our individual target stars, we employed the same iterative reddening approach as described in Harris et al. (2006), assuming zero reddening at 100 pc and full reddening at 250 pc.
3.2 Completeness
To normalize our LF, we employed the standard (
)
method as
introduced by Schmidt (1998). The quantity
describes the maximum volume
in which a star can contribute to the LF:

where,



Table 1: SDSS DR4 spectroscopic hot_std sample incompleteness map.
The
incompleteness factor is not a constant number for all WDs in our
sampled due to the additional isolation criteria placed on hot_stds.
The hot_std incompleteness is a function
of both Galactic latitude, b and object brightness, g. At lower b, the increased
number of stars will make the hot_std unblended criteria harder to meet,
thereby increasing the incompleteness. Similarly, brighter objects cover more
pixels on the CCD, and thus are more likely to be thrown out of the hot_std
category due to the blending or edge criteria.
Eisenstein et al. (2006a) estimates that hot_std incompleteness varies from 0.55 to 0.60 for
30 < b < 40 and from 0.60 to 0.63 when
,
then from 0.40
to 0.55 for g magnitudes
15 < g< 17, and from 0.55 to 0.63 for
mag.
Because we need to know the incompleteness as both b and g vary,
we queried the SDSS database to obtain a full
two-dimensional incompleteness map. We formed a
grid with g varying
from 14-19 in 1 mag bins and b varying from -90 to +90 in 10 degree
bins. Our method was to determine the ratio of hot_stds observed spectroscopically
to the total number of potential hot_stds that passed the photometric
criteria, but not the isolation requirements for each of our (g, b) cells.
The SDSS software pipeline separates blended objects into a parent
(deblended) object and its constituent children. We allowed the parent
object in our search, but not the children.
We also removed all moving objects (e.g., asteroids) and objects too close to the
edge of the CCD chips to make sure that the photometry was properly
determined. Table 1 shows the results of our
calculations. Values in cells are given in
and ``-'' denotes cells
with no
spectroscopic hot_stds in them. Magnitude bins are 1.0 mag wide in g starting
from given magnitudes in the first column. Latitude bins are 10 degrees wide
extending from the values in the first row.
Altogether we found 5778 photometric (PH) hot_standards and
1984 spectroscopically (SP) observed objects formally classified by the
pipeline as a hot_std.
The incompleteness factors of spectroscopic hot_stds for each cell can be
calculated as:

where footprint (

3.3 Results
Using the incompleteness map given in Table 1, we calculated our LF for
bins of 0.5 mag width. The resulting luminosity function
is presented in Fig. 2 along with that of
Harris et al. (2006). It covers an
absolute bolometric magnitude range between -0.75 and 7.25 and provides some
overlap with the cooler LF of Harris et al. (2006).
To determine exactly which bins were affected by our temperature cut, we
calculated the LF for two different cutoff temperatures, our nominal
23 500 K and one at 20 000 K, as discussed in Sect. 2.
The two LFs converge at
,
which we therefore take as the
coolest complete bin in our LF. There appears to be general agreement with
Harris et al. (2006), although there aren't any complete bins in common.
Note the plateau seen in our LF between
(
K).
Eisenstein et al. (2006a)
point out a systematic overestimation of
in
their model fits for the hottest WDs, starting for DAs around
K. At
K, approximately
,
this overestimate is a roughly 10% effect. Furthermore, at these
temperatures, non-LTE (NLTE) effects become important whereas the
Eisenstein et al. (2006a) fits
were based only on LTE models. To correct for these two mis-estimates, we
applied a parabolic temperature correction determined by the Eisenstein et al. (2006a)
comparisons to existing literature fits and NLTE
corrections from Napiwotzki et al. (1999). The resulting LF is shown in
Fig. 3.
As expected, the cool end of this corrected LF remains largely the same, but the
hot end changes considerably, the parabolic temperature adjustments having
the biggest effect.
![]() |
Figure 2:
Our primary WD LFs. The solid line includes all stars in our sample while
the dotted line excludes the more uncertain magnetic and binary WDs.
The two dashed lines show the effect of different temperature cuts: our
preferred cut at 23 500 K and one cut further at 20 000 K to provide more
overlap with the Harris et al. (2006) sample, shown as the solid line with open
circles. The top axis shows the approximate effective
temperature in kilo-Kelvin corresponding to plotted x-axis,
|
Open with DEXTER |
![]() |
Figure 3:
The same as Fig. 2 but with
|
Open with DEXTER |
Table 2: Number of all types of WDs per magnitude bin.
![]() |
Figure 4: Our WD LF as in Fig. 3, but with 1.0 mag bins. |
Open with DEXTER |
With these corrections included, 41 stars' temperatures dropped to below our 23 500 K cutoff. Table 2 lists the LF bin values both before and after these corrections (as shown in Figs. 2 and 3). Columns 2 and 3 list the number of DAs in each bin before and after the corrections. Since the corrections to the DAO, DB, DO, and PG 1159 stars were not needed, they are listed separately in Cols. 4 through 7, respectively. Finally, Col. 8 lists the total number of WDs of all types in each bin.
As Fig. 3 shows, the plateau seen in
Fig. 2 narrows, but remains even after the NLTE and fit
temperature corrections were applied. Figure 4 repeats
the LFs shown in Fig. 2, but with 1.0 mag bins to
increase the signal in each bin; again the plateau remains. While there
are still problems with the hottest DA WDs temperature determination (see
conclusions), a glance at Table 2 shows that the DAs start appearing at
,
precisely where our plateau begins. It appears we are
seeing the emergence of the non-DAs into the LF.
![]() |
Figure 5:
DA (solid circles) and non-DA (open squares) LFs plotted with the PG
survey DA LF (open circles) of Liebert et al. (2005).
The magnitude bin width is 1.0 and the bins are centered on integer
magnitudes. Thinner line right below SDSS DA LF is the DA LF without magnetic
DAs and binary DA+M systems. In case of non-DA LF removing magnetic and binary
DB+M systems from the WD sample affects the LF near
|
Open with DEXTER |
4 DA and non-DA luminosity functions
Table 2 shows that the hot end of our LF is made up entirely of
non-DA WDs. We therefore divided our main LF into two LFs, one with only
DAs and one for non-DAs. We kept the 23 500 K
cutoff for the
DA-only LF, but were able to decrease the temperature cutoff to 16 000 K for
the non-DA LF. Using bins of 1.0 mag centered on half integer and
integer magnitudes, we plot the resulting
new LFs in Figs. 5 and 6, respectively.
We also compare to the DA LF of Liebert et al. (2005), after converting it from
MV into
space densities.
![]() |
Figure 6: The same as in Fig. 5 but bins are centered on half integer magnitudes. In both plots, we can see the appearance of the first DAs affect a dip in the non-DA LF. |
Open with DEXTER |
The first thing to notice is that our new DA LF and that of Liebert et al. (2005)
are in very good agreement.
The hot end of the non-DA LF has an interesting shape, with a sudden drop
in the LF near
.
,
however, this is precisely where the
DAs start appearing. If we follow the initial slope of the non-DA LF and
extend it before the drop-off to
,
it leads directly to the
start of the DA LF. Thus, it seems this drop off is simply due to the
appearance of DAs. These luminosity functions represent the first time this
separation of the WD population has been seen. Table 3 provides the numbers
behind the separate DA and non-DA LFs as shown in
Fig. 6. The last DA and non-DA bins
(marked with asterisks) are incomplete due to our temperature cutoffs at
23 500 K and 16 000 K respectively.
Table 3: DA and non-DA WD distribution per magnitude bin.
5 Conclusions
As shown, our WD LF agrees very well with previous efforts, but extends to
significantly hotter stars and shows clear
emergence of DAs and non-DAs in the general WD cooling curve. However, due
to the problems with hot DA WD temperature determination the DA and non-DA
space densities cannot be compared directly at the hot end of these LFs.
Although the systematic DA
overestimation of Eisenstein et al. (2006a) WD
catalog data has
been corrected and the pure hydrogen NLTE vectors have been applied, yet
there are more severe problems affecting the hottest DA WD temperature and
determination. These include lower spectra quality of the
fainter HSs or the presence of metals in the hot DA WD atmosphere
visible as the onset of the
Balmer line problem (Vennes 1999) in the spectra. Barstow et al. (2003) and
Vennes et al. (2005) investigated the effect of metallicity influence on the hot DA WD
(above 50 000 K) temperature determination. They have shown that, due to
heavy elements present in the DA atmosphere, in some cases the hot DA
temperatures can be overestimated by at least
,
which translates
to
1
more luminous WDs. In such circumstances the space density of the
DA WDs seems reasonably well established for
mag and
is not certain for the more luminous part of the DA LF.
In that regard, the non-DA LF is likely the most reliable out of both our
LFs. It covers -0.5 to 9.5 bolometric magnitude range and shows the drop
off near
which we interpret as the non-DA to DA atmosphere
conversion (Figs. 5, 6).
We can see that in the DB gap temperature range, 30 000-43 000 K
(
), our non-DA LF has only a slight space density decrease,
suggesting the number of recently discovered hot DB stars (Eisenstein et al. 2006a,b)
has nearly, if not completely, filled the ``gap''.
Note that the error bars presented in the figures are formal errors, calculated as a sum
of
for each luminosity bin, and they do not take observed or
absolute bolometric magnitude error determinations into account. In fact the
method underestimates LF error bars by a factor of 10 (Geijo et al. 2006) near
the 9-7
magnitude range where our LF has its cool end, and in that regard
the agreement between our LF and others can have larger tolerance. In their
theoretical work Geijo et al. (2006) did not make any LF error estimations for the hotter
regions our LF covers, so we cannot determine the error quality in that range,
but we can expect the ones we present are also underestimated.
In short, we have calculated a new WD LF with an extended range of luminosities based on the well-understood and explored SDSS WD data set. However, to improve the quality of the LF at its hot end, new follow up observations of low signal to noise spectra are necessary as well as careful analysis of the metal content in the WD spectra. Further exploration of our LF should contribute a number of new insights into early white dwarf evolution.
Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/.
The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh,University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.
References
- Abazajian, K., Adelman-McCarthy, J. K., Agüeros, M. A., et al. 2003, AJ, 126, 2081 [NASA ADS] [CrossRef]
- Abazajian, K., Adelman-McCarthy, J. K., Agüeros, M. A., et al. 2004, AJ, 128, 502 [NASA ADS] [CrossRef]
- Abazajian, K., Adelman-McCarthy, J. K., Agüeros, M. A., et al. 2005, AJ, 129, 1755 [NASA ADS] [CrossRef]
- Adelman-McCarthy, J. K., Agüeros, M. A., Allam, S. S., et al. 2006, ApJS, 162, 38 [NASA ADS] [CrossRef]
- Barstow, M. A., Good, S. A., Burleigh, M. R., et al. 2003, MNRAS, 344, 562 [NASA ADS] [CrossRef]
- Blanton, M. R., Lin, H., Lupton, R. H., et al. 2003, AJ, 125, 2276 [NASA ADS] [CrossRef]
- Bloecker, T. 1995, A&A, 299, 755 [NASA ADS]
- DeGennaro, S., Hippel, T., Winget, D. E., et al. 2008, AJ, 135, 1 [NASA ADS] [CrossRef]
- Eisenstein, D. J., Liebert, J., Harris, H. C., et al. 2006a, ApJ, 167, 40 [NASA ADS] [CrossRef]
- Eisenstein, D. J., Liebert, J., Koester, D., et al. 2006b, AJ, 132, 676 [NASA ADS] [CrossRef]
- Evans, D. W. 1992, MNRAS, 255, 521 [NASA ADS]
- Finley, D. S., Koester, D., & Basri, G. 1997, ApJ, 488, 375 [NASA ADS] [CrossRef]
- Fleming, T. A., Liebert, J., & Green, R. F. 1986, ApJ, 308, 176 [NASA ADS] [CrossRef]
- Fukugita, M., Ichikawa, T., Gunn, J.E., et al. 1996, AJ, 111, 1748 [NASA ADS] [CrossRef]
- Geijo, E. M., Torres, S., Isern, J., & Garcia-Berro, E., MNRAS, in preparation
- Girardi, L., Bertelli, G., Bressan, A., et al. 2002, A&A, 391, 195 [NASA ADS] [CrossRef] [EDP Sciences]
- Green, R. F., Schmidt, M., & Liebert, J. 1986, ApJS, 61, 305 [NASA ADS] [CrossRef]
- Gunn, J. E., Carr, M., Rockosi, C., et al. 1998, AJ, 116, 3040 [NASA ADS] [CrossRef]
- Gunn, J. E., Siegmund, W. A., Mannery, E. J., et al. 2006, AJ, 131, 2332 [NASA ADS] [CrossRef]
- Harris, H. C., Munn, J. A., Kilic, M., et al. 2006, AJ, 131, 571 [NASA ADS] [CrossRef]
- Hogg, D. W., Finkbeiner, D. P., Schlegel, D. J., & Gunn, J. E. 2001, AJ, 122, 2129 [NASA ADS] [CrossRef]
- Hu, Q., Wu, C., & Wu, X.-B. 2007, A&A, 466, 627 [NASA ADS] [CrossRef] [EDP Sciences]
- Hügelmeyer, S. D., Dreizler, S., Homeier, D., et al. 2006, A&A, 454, 617 [NASA ADS] [CrossRef] [EDP Sciences]
- Isern, J., Garcia-Berro, E., Torres, S., & Catalan, S. 2008, ApJ, 682, L109 [NASA ADS] [CrossRef]
- Ivezic, Z., Lupton, R. H., Schlegel, D., et al. 2004, Astron. Nachrichten, 325, 583 [NASA ADS] [CrossRef]
- Kleinman, S. J., Harris, H. C., Eisenstein, D. J., et al. 2004, AJ, 607, 426 [NASA ADS] [CrossRef]
- Krzesinski, J., Nitta, A., Kleinman, S. J., et al. 2004, A&A, 417, 1093 [NASA ADS] [CrossRef] [EDP Sciences]
- Liebert, J., Dahn, C. C., & Monet, D. G. 1988, ApJ, 332, 891 [NASA ADS] [CrossRef]
- Liebert, J., Bergeron, P., & Holberg, J. B. 2005, ApJS, 156, 47 [NASA ADS] [CrossRef]
- Monet, D. G., Levine, S. E., Canzian, B., et al. 2003, AJ, 125, 984 [NASA ADS] [CrossRef]
- Napiwotzki, R., Green, P. J., & Saffer, R. A. 1999, ApJ, 517, 399 [NASA ADS] [CrossRef]
- O'Brien, M. S., & Kawaler, S. D. 2000, ApJ, 539, 372 [NASA ADS] [CrossRef]
- Oke, J. B., & Gunn, J. E. 1983, ApJ, 266, 713 [NASA ADS] [CrossRef]
- Oswalt, T. D., Smith, J. A., Wood, M. A., & Hintzen, P. 1996, Nature, 382, 692 [NASA ADS] [CrossRef]
- Pier, J. R., Munn, J. A., Hindsley, R. B., et al. 2003, AJ, 125, 1559 [NASA ADS] [CrossRef]
- Rauch, T., & Deetjen, J. L. 2003, Stellar Atmosphere Modeling, 288, 103 [NASA ADS]
- Schmidt, M. 1968, ApJ, 151, 393 [NASA ADS] [CrossRef]
- Schoenberner, D. 1983, ApJ, 272, 708 [NASA ADS] [CrossRef]
- Smith, J. A., Tucker, D. L., Kent, S., et al. 2002, AJ, 123, 2121 [NASA ADS] [CrossRef]
- Stoughton, C., Lupton, R. H., Bernardi, M., et al. 2002, AJ, 123, 485 [NASA ADS] [CrossRef]
- Tucker, D., Kent, S., Richmond, M. W., et al. 2006, Astron. Nachrichten, 327, 821 [NASA ADS] [CrossRef]
- Vennes, S. 1999, ApJ, 525, 995 [NASA ADS] [CrossRef]
- Vennes, S., Smith, R. J., Boyle, B. J., et al. 2002, MNRAS, 335, 673 [NASA ADS] [CrossRef]
- Vennes, S., Chayer, P., Dupuis, J., & Lanz, T. 2005, ASPC, 334, 1985
- Wesemael, F., Green, R. F., & Liebert, J. 1985, ApJS, 307, 65
- Wood, M. A. 1995, in White Dwarfs, Proc. of the 9th European Workshop on White Dwarfs ed. D. Koester, & K. Werner (Berlin, Heidelberg, New York: Springer-Verlag), Lecture Notes in Physics, 443, 41
- Wood, P. R., & Faulkner, D. J. 1986, ApJ, 307, 65 [CrossRef]
- York, D. G., Adelman, J., Anderson, J. E., Jr., et al. 2000, AJ, 120, 1579 [NASA ADS] [CrossRef]
All Tables
Table 1: SDSS DR4 spectroscopic hot_std sample incompleteness map.
Table 2: Number of all types of WDs per magnitude bin.
Table 3: DA and non-DA WD distribution per magnitude bin.
All Figures
![]() |
Figure 1:
The photometric
|
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Our primary WD LFs. The solid line includes all stars in our sample while
the dotted line excludes the more uncertain magnetic and binary WDs.
The two dashed lines show the effect of different temperature cuts: our
preferred cut at 23 500 K and one cut further at 20 000 K to provide more
overlap with the Harris et al. (2006) sample, shown as the solid line with open
circles. The top axis shows the approximate effective
temperature in kilo-Kelvin corresponding to plotted x-axis,
|
Open with DEXTER | |
In the text |
![]() |
Figure 3:
The same as Fig. 2 but with
|
Open with DEXTER | |
In the text |
![]() |
Figure 4: Our WD LF as in Fig. 3, but with 1.0 mag bins. |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
DA (solid circles) and non-DA (open squares) LFs plotted with the PG
survey DA LF (open circles) of Liebert et al. (2005).
The magnitude bin width is 1.0 and the bins are centered on integer
magnitudes. Thinner line right below SDSS DA LF is the DA LF without magnetic
DAs and binary DA+M systems. In case of non-DA LF removing magnetic and binary
DB+M systems from the WD sample affects the LF near
|
Open with DEXTER | |
In the text |
![]() |
Figure 6: The same as in Fig. 5 but bins are centered on half integer magnitudes. In both plots, we can see the appearance of the first DAs affect a dip in the non-DA LF. |
Open with DEXTER | |
In the text |
Copyright ESO 2009
Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.
Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.
Initial download of the metrics may take a while.