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1997MNRAS.286..369M
Mon. Not. R. Astron. Soc. 286,
369-383 (1997)
An EUV-selected sample of DA white dwarfs from the ROSAT All-Sky
Survey - I. Optically derived stellar parameters
M. C. Marsh/*t M. A. Barstow/* D. A. Buckley,2 M. R. Burleigh/*
J. B. Holberg,3:j: D. Koester,4 D. Q'Donoghue,5 A. J. Penny6 t and
A. E. Sansom7t
lDepartment of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH
South African Astronomical Observatory, Cape Town, South Africa
3Lunar and Planetary Laboratory, University of Arizona, Gould-Simpson Building, Tucson, AZ 85721, USA
41nstitut jUr Astronomie und Astrophysik, Universitiit Kie~ D-24098, Kie~ Germany
5Department of Astronomy, University of Cape Town, Rondebosch 7700, South Africa
6Astrophysics Division, Rutherford Appleton Laboratory, Didcot, Oxfordshire OXll OQX
7 Centre for Astrophysics, University of Central Lancashire, Preston PR1 2HE
2
1996 November 7. Received 1996 November 5; in original form 1996 May 30
ABSTRACT
One of the most important results of the ROSAT All-Sky X-ray and EUY Surveys
has been the detection of a large number of white dwarfs, allowing a detailed study
of the general properties of DA white dwarf atmospheres to be carried out.
However, this work relies on a knowledge of the effective temperatures, surface
gravities and visual magnitudes. We present analyses of optical data from our
follow-up programme, including the values of T eff , logg, mv and stellar mass
obtained. We also list the PSPC and WFC count rates for each star. The distribution
of masses and surface gravities derived from the optical work also provides
important information about the sample of white dwarfs and gives some indication
of possible selection biases. Comparing the ROSAT sample of 89 stars with the
cooler optically selected sample of Bergeron, Saffer & Liebert reveals a statistically
significant excess (~3 times the number found by them) of hot, massive DAs, which
may represent a population of coalesced binary white dwarfs. As the optical and
EUY samples do not cover the same range of white dwarf temperatures, the highmass excess may partly arise from differences in the cooling rates of 'normal'
(~O.6-M0) and massive (> 1.0-M0) stars. Consequently, this feature is not
necessarily the result of selection on the basis of EUY flux, and it might also be
present in an optically selected sample covering an appropriate temperature
range.
Key words: stars: atmospheres - stars: luminosity function, mass function - white
dwarfs.
1 INTRODUCTION
Among the most significant problems in the study of hot
white dwarf evolution has been the existence of two distinct
groups having either H- or He-dominated atmospheres and
the possible relationships between these groups and their
*Guest Observer with the SAAO l.9-m telescope.
tGuest Observer with the l.O-m Jacobus Kapteyn Telescope of the
ING.
:j:Guest Observer with the Steward Observatory 2.3-m telescope.
proposed progenitors, the diverse types of central stars of
planetary nebulae (CPN). While the very hottest H-rich DA
white dwarfs outnumber He-rich DOs by a factor of 7
(Fleming, Liebert & Green 1986), the relative number of Hand He-rich CPN is only about 3:1. In addition, there is an
apparent absence of He-rich stars in the temperature range
30000-45000 K, the so-called DO-DB gap, suggesting that
H- and He-dominated groups are not entirely distinct.
Several competing processes can affect the composition of a
white dwarf atmosphere. He and heavier elements tend to
sink out of the photosphere under the influence of gravity,
©1997 RAS
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
Accepted
1997MNRAS.286..369M
370 M C. Marsh et al.
results with the white dwarf evolutionary sequences of
Wood (1992), they were able to estimate the stellar masses
and, consequently, obtain the mass distribution of the
sample. This spectroscopic technique provides a more
detailed picure than earlier photometric studies. However,
statistical work of this nature often has a number of limitations. The BSL sample of 129 DA white dwarfs is the most
extensive study reported in detail, for which the measured
mean mass was 0.56 Mo. However, the sample spanned only
the temperature range from ~ 13 000 to 40000 K, and the
number of white dwarfs with temperatures above
~25 000 K, the point at which X-ray/EUV emission
becomes significant, made up no more than 25 per cent of
the total. Consequently, the BSL sample does not overlap
significantly with the stars in the ROSAT survey; indeed,
only eight of their stars have been detected as X-ray/EUV
. sources, and their results are not necessarily representative
..of the hotter, younger population we have observed.
More recently, some results have been presented extending study of the mass distribution up to higher temperatures. Liebert & Bergeron (1995) have studied a sample of
'white dwarfs with temperatures above 15 000 K, drawn from
the Palomar Green survey and including 200 white dwarfs.
They obtain a slightly higher mean mass of 0.582 when
carrying out the analysis with the earlier evolutionary
models, which did not include an outer hydrogen layer.
. However, using new models with thick hydrogen layers
(10- 4 Mo; Wood 1995), which are more in keeping with
theoretical predictions, yields a still higher value of 0.609
(note: wid). these models the BSL sample mean increases to
0.590). Finley (1995) reports a study of 177 DA white dwarfs
with temperatures above 24000 K, obtaining mean masses
of 0.601 and 0.63 Mo (Finley, private communication) with
the zero and thick hydrogen layer mass models respectively,
higher than either BSL or Liebert & Bergeron (1995). However, this sample is rather more heterogeneous than that
drawn from the PG survey, containing a significant fraction
of EUV-selected objects besides those discovered by optical
means.
, .
All these re~rled mass distributions are dominated
by stars selected by optical observations. As Liebert &
Bergeron (1995) point out, these are strongly biased against
the highest· mass stars with the smallest radii. Selecting stars
on the basis' of their EUVemission may diminish this effect
to some extent. Radiative levitation calculations (e.g., most
recently, Chayer et al. 1994, 1995) predict that photospheric
abundances of heavy elements should depend strongly on
gravity, higher gravity stars having a lower heavy-element
content. Consequently, those white dwarfs with the greatest
mass should also be the most luminous in the EUV. How
much this can counteract the effect of smaller radius will,
however, depend on the stellar temperature, since the
radiative levitation mechanism becomes efficient only above
40000 K In addition, other, different, selection effects may
operate on the EUV sample, for example, interstellar
absorption and effective temperature, to which the
observed EUV flux is very sensitive. We present here a
detailed comparison of the EUV-selected sample of white
dwarfs with those drawn from the optical catalogues. A
larger number of low-gravity objects are found in the
ROSAT WFC survey, which is consistent with the bias to
high white dwarf temperatures. However, more import© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
but this can be counteracted by radiation pressure. Convective mixing, accretion or mass loss via a weak wind may
also playa Significant role in determining the heavy-element
content of white dwarf atmospheres.
Shipman (1976) made the now universally accepted
proposition that the X-ray/EUV flux from DA white dwarfs
was, in fact, from their hot photospheres. He pointed out
that, since the emission· would emerge from the· hotter,
deeper layers of the stellar atmosphere, X-ray/EUV observations should provide a sensitive probe of its structure and
composition. This important idea has underpinned all subsequent X-ray/EUV studies of white dwarfs. In particular, it
was possible to demonstrate, from an analysis of 30 catalogued DA white dwarfs detected in the ROSAT All-Sky Xray and EUV Surveys, that most white dwarfs hotter than
~ 40000 K contain significant quantities of elements
heavier than either H or He, while below this temperature
the stars can be represented by more or less pure H atmospheres (Barstow et al. 1993). Subsequent studies, increasing the numbers of white dwarfs in the sample (e.g. Marsh et
al. 1995; Wolff, Jordan & Koester 1996), have confirmed
these results.
The data provided by ROSAT have led to a clear breakthrough in our understanding of the general properties of
the atmospheres of the hot DA white dwarf population.
However, the full exploitation of this resource would not
have been possible without prior knowledge of the temperatures, gravities and visual magnitudes of the stars detected.
These parameters, derived for the most part from optical
observations, serve to specify the EUV/X-ray flux for a
given model composition. In principle, if the composition of
the white dwarf atmosphere is known, the entire spectrum
can be modelled from the temperature and gravity. The
visual magnitude then acts as a convenient normalization
factor defining the ratio between stellar distance and radius.
Thus the remaining free parameters, related to the photospheric and intervening interstellar opacity, can be determined from the data.
. When ROSAT was launched, this necessary information
was available for only a relativ:ely small number of already
catalogued white dwarfs, and this was the main factor limiting the size of the sample in the first analysis of the survey.
Approximately 176 white dwarfs Were ultimately detected
by the X-ray telescope (Fleming·et aL 1996), and 120 were
seen in the EUV. CQnseq~ent1y, the 'original survey sample
can potentially be increased by a factor of 4, although in
practice not all ROSAT detections are bright enough to
provide useful information. These additional objects
include a number of known white dwarfs not previously
reported as X-ray sources, and not well studied at optical
wavelengths, together with a substantial number of newly
discovered stars. This latter group account for ~ 50 per cent
of the total number of white dwarfs in the ROSAT catalogues.
A determination of effective temperature and surface
gravity of a white dwarf can be obtained from the profiles of
the hydrogen Balmer lines in the optical spectrum, comparing observations with the calculations of theoretical model
atmospheres. This technique was pioneered by Holberg et
al. (1985) and Kidder (1991). Bergeron, Saffer & Liebert
(1992, hereafter BSL), have used it to carry out a systematic
study of the DA white dwarf population. Combining these
1997MNRAS.286..369M
A UV-selected sample of DA white dwarfs
antly, we find an excess of high-mass objects in the ROSAT
group. This might arise from the preferential selection of
those white dwarfs with the highest EUY luminosity, with
high gravity counteracting the effect of radiative levitation,
lowering the abundances of heavy elements and, as a consequence, the photospheric opacity in the EUY. Alternatively,
the excess may be a product of the bias to higher white
dwarf temperatures. As Finley (1995) points out, there is a
sharp decrease in the cooling rate of massive (> 1 Mo)
white dwarfs in the range 40000-50000 K, which would
increase the numbers of such stars in this interval, compared
to the 'normal' (lower mass) population.
2 THE ROSAT SKY SURVEY AND OPTICAL
IDENTIFICATION PROGRAMME
3 OPTICAL SPECTROSCOPY
3.1
Observations
As hot DA white dwarfs are thermal sources of radiation
with spectra peaking in the EUY and soft X-ray, their spectral shape and hence their observed count rates in the
ROSAT bandpasses are very sensitive to T eff . It is therefore
important that Teff is well determined, in order to fully
utilize the ROSAT survey data. This work demands relatively high signal-to-noise spectra (SIN;:::: 50) of moderate
resolution (~8 A), and the data from the optical identification programme were, in general, unable to fulfil this
requirement. Consequently, a programme of spectroscopic
observations was undertaken in both northern and southern
hemispheres to obtain data for the newly discovered group
of white dwarfs and those catalogued stars where suitable
spectra did not exist.
Southern hemisphere observations were made in 1993
October and 1994 March usirig the 1.9-m Radcliffe telescope of the South African Astronomical Observatory
(SAAO), with the Unit spectrograph and Reticon photon
counting system (RPCS). The RPCS has two arrays, one
which accumulates energy from the source while the other
records the sky background through an adjacent aperture.
The grating was blazed to cover a wavelength range of
3600-5100 A and had a reciprocal dispersion of 75 A mm -1,
yielding a resolution of ;::;2.5 A (FWHM). Flat-field spectra
were obtained at the start and end of each night, and wavelength calibration was provided by a CuAr arc lamp which
was observed at the beginning and the end of each spectrum
acquisition. Several spectrophotometric standards such as
LDS 678A were observed in order to convert the observed
counts to flux.
Stars in the northern hemisphere were observed with the
Steward Observatory 2.3-m telescope on Kitt Peak. This
instrument is equipped with a Boller & Chivens spectrograph and a UY-flooded 800 x 1200 Loral CCD detector.
The grating was also blazed to cover the spectral range
occupied by the Balmer lines, and gave a resolution of ;::; 8 A
(FWHM). The spectra were then flat-fielded and converted
to flux using observations of standard stars. A detailed
description of thee reduction procedures can be found in
Kidder (1991). All data were acquired at the parallactic
angle using 1.25- or 2.5-arcsec slits, depending upon the
seeing conditions. Fig. 1 shows a representative sample of
18 flux-calibrated spectra obtained in both northern and
southern hemisphere programmes.
3.2 Determination of temperature and gravity
Comparison of previously published DA temperatures
derived from Balmer line fitting reveals that systematic differences in the values determined may occur as a result of
using different model spectra. This problem was first noted
by Bergeron, Wesemael & Fontaine (1991), who identified
the cause as arising from the treatment of the high-lying
levels of hydrogen in the model atmosphere calculations.
For example, work included in Kidder (1991) used an
extended grid of line-blanketed homogeneous LTE models
computed by F. Wesemael (Wesemael et al. 1980), whereas
Finley, Koester & Basri (1997) employed fully line-blanketed homogeneous LTE models developed by Koester
(see Koester 1991 and Finley et a1. 1997). A comparison of
these results shows that, while the temperatures obtained by
Kidder (1991) are similar to those of Finley et aI., the surface gravities of the former sample are on average 0.08 dex
higher (Finley, private communication). The more recent
Koester models do take into account higher excitation levels
of the H atoms, giving more reliable results, and these are
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
The ROSAT satellite, launched in 1990 June, carries two 00aligned imaging instruments, an X-ray telescope (XRT)
(Pfefferman et al. 1986) and the Wide Field Camera (WFC)
(Sims et al. 1990), covering the EUY waveband. The prime
objective of the mission was to carry out the first survey of
the entire sky in both the EUY and X-ray during the first 6
months of operation. The EUY survey was conducted in
two bands, the S1 filter covering the range 60-140A
(90-200 eV), and the S2 filter for 112-200 A (60-110 eV).
A complete catalogue of X-ray sources has not yet
appeared, but the white dwarf fraction has been published
by Fleming et al. (1996), who list the 'C band' (0.1-0.28
ke V) count rates for each detection. Two catalogues of
EUY sources have been produced from analysis of the
WFC data. The Bright Source Catalogue (BSC) (Pounds et
a1. 1993) contains 119 identified white dwarfs, while a
second catalogue (2RE) with lower detection thresholds
yields only six more (Pye et a1. 1995), although the total
number of sources is increased from 383 to 479. Around 35
2RE EUY sources have yet to be identified optically but, as
the vast majority (;::;75 per cent) of the identified new
sources are active stars, we may expect only a few further
white dwarf discoveries. Hence the EUY-selected sample of
white dwarfs is drawn predominantly from the BSC.
Once EUY sources had been detected in the initial survey
analyses, the next step was to identify the optical counterpart, initially through cross-correlation of source positions
with astronomical catalogues. When no catalogued counterpart could be found, the EUY source fields were included in
a programme of optical spectroscopic follow-up observations (Mason et a1. 1995). From this work ;::;50 new white
dwarfs were found. The complete sample is listed in Table 1,
ordered by BSC name and including the most commonly
used name for those stars already catalogued. Also indicated are a number of objects found in astronomical catalogues but which were not previously known to be white
dwarfs (e.g., REJ0237 -12=PHL 1400).
371
1997MNRAS.286..369M
372 M, C. Marsh et ai,
Table 1. Epoch-corrected count rates from the ROSAT All-Sky Survey,
Alternative
Name
new ID
GD 2
MCT
GD 659
GD 683
GD 984
PG 0136 + 251
GD 1401
GD 421
LB 1628
Feige 24
PHL 1400
LB 1663
GD 50
V471 TAU
MCT
G191 B2B
new ID
new ID
GD 257
new ID
GD 71
new ID
new ID
new ID
SIRIUS B
GD 80
new ID
new ID
PG 0824 + 289
new ID
new ID
new ID
PG 0904 +511
new ID
new ID
PG 1026 +454
new ID
EG 070
GD 123
new ID
PG 1041 +580
new ID
PG 1057 + 719
TON 61
PG 1123 + 189
PG 1125 - 026
PG 1145 + 188
PG 1234 +482
EG 187
HZ 43
new ID
GD 336
new ID
new ID
CD - 3810980
KUV 433-3
new ID
PG 1725 +587
new ID
KUV 18004+6
new ID
BPM 93487
PSPC"
356.0
2520.0
3790.0
3004.0
87.0
1745.0
632.0
32.0
158.0
50.0
14.0
329.0
230.0
2086.0
1110.0
106.0
40.0
971.0
190.0
725.0
373.0
2800.0
4.0
132.0
9.0
1506.0
19730.0
682.0
3536.0
455.0
1833.0
714.0
32.0
2154.0
81.0
61.0
4744.0
286.0
288.0
196.0
9068.0
242.0
294.0
1634.0
223.0
96.0
1192.0
508.0
436.0
50.0
88.0
2050.0
7640.0
37170.0
45.0
153.0
153.0
225.0
969.0
468.0
157.0
1387.0
24.0
1248.0
812.0
1691.0
1423.0
Survey Count Rates (c ks- 1 )
1"."
Sl
1".
S2
47.0
37.3
22.5
6.8
121.0
148.9 12.4
50.4
190.0
625.7 35.9 1008.2
30.0
765.5 30.1 2079.1
19.0
43.5 12.1
134.2
88.0
259.4 14.7 656.3
63.0
52.2
6.8
50.6
13.5
24.7
11.0
5.6
20.0
23.3
4.1
49.4
16.0
24.8
4.5
135.1
758.0
4.9
2.3
39.0
55.6 10.2
121.6
68.9
76.0
6.8
210.4
71.0
291.7 16.2 500.7
250.0
340.0 20.0 1040.0
27.0
120.5 10.2 2567.3
62.2
22.0
8.3 2824.1
168.3 12.3
282.1
64.0
34.0
55.7
7.2
205.7
100,4 8.4
13.0
94.4
23.0
41.0
4.2
15.4
113.0
679.0 21.9 2281.3
7.4
3.2
34.0
27.0
10.7
4.3
43.6
6.4
3.2
428.0
66.0
142.0
9.0
255.0
290.0 3200.2 41.8 8187.4
55.0
106.7 8.6
279.8
1127.0
539.3 15.9 867.8
109.5
45.0
9.9
369.9
80:0
297.9 15.3 987.6
53.0
46.2
7.7
52.8
15.0
20.2
5.9
64.2
238.2 16.8 448.2
102.0
23.0
26.9
5.6
37.4
19.0
35.3
47.6
528:0 19.9
116.0
594.4
28.3
27.0
5.9
77.0
30.0
58.6
8.4
137.7
22.0
31.4
4.5
88.7
151.0 1136.5 25.8 2689.8
27.0
35.5
118.0
7.1
57.2
6.7 212.8
26.0
69.0
226.6 11.2 477.9
32.0
49.2
6.7
121.6
16.0
43.1
7.7
128.3
74.0
120.3 19.9
325.0
42.0
90.1 11.7
88.8
183.2 13.0 990.5
33.0
22.7 4.8
42.2
15.0
41.6
100.8
7.1
18.0
30.0
515.6 20.1 1890.2
163.0 1509.5 37.6 4749.8
412.0 12157.6 92.9 36982.6
10.0
10.6 4.7
15.4
22.6
4.5
20.4
25.0
26.0
17.0
3.4
38.7
17.3 7.1
24.0
23.2
54.0
199.5 14.3 605.8
29.0
51.7
7.1
51.4
17.3
36.0
1.0
33.3
28.0
47.7
5.7
14.0
10.0
11.4
1.1
11.8
106.2 18.9
306.0
267.6
237.0
51.1
2.3
21.5
104.5
3.4 112.0
112.0
260.7 15.0 697.6
75.0
1".
9.7
40.3
44.6
15.7
23.6
7.5
7.5
5.0
9.7
30.1
12.9
11.8
19.1
29.0
42.5
40.8
14:2
14.3
9.2
6.2
37.9
5.1
6.2
15.6
15.0
68.7
15.7
23.7
19.0
31.6
10.6
7.8
22.4
8.5
8.5
24.1
8.8
12.2
10.7
37.4
14.3
12.8
18.2
9.6
11.7
20.1
13.2
34.2
7.8
10.0
40.8
54.1
148.4
6.6
6.4
6.4
9.1
22.3
7.1
8.1
1.1
26.8
2.2
4.3
23.9
Mean Julian
Date
48179.7
48178.8
48195.9
48223.8
48226.8
48178.8
48182.6
48178.8
48111.1
48178.8
48183.3
48182.6
48179.7
48108.6
48110.6
48115.5
48125.0
48118.6
48122.1
48129.9
48129.5
48129.5
48133.9
48141.8
48146.2
48141.8
48146.2
48149.2
48183.3
4;8159.5
'48159.5
48162.4
. 48215.2
. :4817'3.3
48174.5
48163.4
48204.3
48208.3
48218.9
48175.4
48174.0
48210.8
48175.4
48175.4
48173.0
48165.8
48164.4
48204.3
48210.8
48219.4
48215.7
48211.3
48165.8
48165.8
48208.8
48179.6
48180.6
48113.0
48116.7
48111.6
48114.0
48183.3
48183.3
48133.5
48183.3
48183.3
48146.2
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
WFC
Name
REJ0003 +43
REJ0007 + 33
REJ0029 - 63
REJ0053 - 32
REJ0108 - 35
REJ0134 -16
REJ0138 +25
REJ0148 - 25
REJ0151 + 67
REJ0230 - 47
REJ0235 + 03"
REJ0237 -12
REJ0322 - 53
REJ0348 - 00
REJ0350 + 17 6
REJ0457 - 28
REJ0505 + 52
REJ0512 - 41
REJ0521-1O
REJ0550 + 00
REJ0550 - 24
REJQ552 + 15
REJ0558 - 37
REJ0605 - 48
REJ0623 - 37
REJ0632 - 05'
REJ0645 -16
REJ0654 - 02
REJ0715 -70
REJ0720 - 31"
REJ0723 - 27
REJ0827 + 28
REJ0831- 53
REJ0841 + 03
REJ0902 - 04
REJ0907 + 50
REJ1016 - 05"
REJ1019 -14
REJ1024 - 30
REJ1029 +45
REJ1032 + 53
REJ1033 -11
REJ1036 +46
REJ1043 +49
REJI044 + 57
REJ 1058 - 38
REJ1100 + 71
REJ1112 +24
REJ1126 + 18
REJ1128 - 02
REJ1148 + 18
REJI236 +47
REJ1257 + 22
REJI316 +29
REJ1340 +60
REJ1431 + 37
REJ1529 +48
REJ1614 - 08
REJ1623 - 39
REJ1638 + 35
REJ1650 +40
REJ1726 +58
REJ1738 +66
REJ1746 -70
REJ1800 +68
REJ1820 +58
REJ1847 + 01
1997MNRAS.286..369M
A UV-selected sample of DA white dwarfs
373
Table 1 - continued
WFC
Name
new ID
new ID
AOO 2000 - 56
new ID
L 210 -114
MCT
new ID
GD 391
GD 394
new ID
new ID
MCT
new ID
new ID
new ID
PHL 0396
GD 246
new ID
MCT
PHL 580
PSPC"
Survey Count Rates (c ks- 1 )
1"."
S2
Sl
1".
316.0
46.0
404.0
82.0
88.0
27.0
12040.0 430.0
935.0
34.0
200.0
55.0
382.0
39.0
442.0
30.0
9.0
71.0
65.0
1282.0
464.0
48.0
158.0
33.0
1650.0
84.0
6900.0 175.0
223.0
21.0
220.0
38.0
13.0
87.0
20.0
6467.0 114.0
1049.0
82.0
55.0
22.0
25.0
136.0
36.3
19.6
22.8
1274.8
61.8
26.1
61.2
22.3
23.1
418.1
57.1
32.5
166.5
1146.3
18.9
75.3
15.6
39.4
1613.6
261.3
40.6
36.9
8.5
4.6
10.8
45.6
6.7
12.0
9.8
5.6
5.8
13.2
7.8
6.7
13.4
29.0
4.7
12.4
6.7
16.2
35.9
20.7
18.5
9.5
48.8
23.9
41.9
2626.6
33.2
32.5
181.0
18.1
27.2
1411.3
98.2
85.5
181.3
2673.7
75.0
133.8
1104.8
70.9
4711.3
656.8
364.1
64.2
1".
11.4
5.4
12.6
72.3 .
6.4
1.0.~
18·.9
7.6
23.4
lO.7
10.7
12.8
43.7
11.0
12.8
30.9
14.2
57.8
29.3
27.3
12.2
Mean Julian
Date
48142.3
48192.7
48153.4
48153.4
48176.4
48155.2
. 48157.7
48172.0
48195.9
48220.4
48173.5
48174.5
48172.5
48171.3
48210.8
48176.4
48172.5
48197.4
48219.4
48192.7
48197.4
48215.2
Notes
*Count S-l in the 0.1-0.28 keY band.
'DAD stars.
'SI rate is PV/Cal data.
bRates are after subtraction of K-star component and removal of eclipses (see Barstow et al.
1992).
CWFC rates derived from 2RE catalogue (Pye et al. 1995).
dpSPC rate is after subtraction of a dMe companion, and should be treated as an upper limit
(see Barstow et al. 1995a).
'-' in the error column indicates that the count rate is an upper limit.
used in the current work. Consequently, as part of our programme, we have reanalysed 48 spectra for which Teff and
logg have already been published (e.g. Kidder 1991; Barstow et al. 1993) to ensure that we have an internally consistent set of values for the subsequent detailed study of this
white dwarf sample.
DA white dwarfs are defined as having no detectable
helium in their optical spectrum. However, a potentially
important effect that may influence the results of Balmer
line fitting is discussed by Bergeron et al. (1994). They have
shown that for DA stars, homogeneous helium abundances
of He!H ~ 10-5 -10- 4 , which yield no detectable He II
absorption at either 4686 or 1640 A, can significantly reduce
the derived effective temperature compared to that for a
pure hydrogen atmosphere. For example, a 52000 K DA
with a homogeneous helium abundance of He!H=10- 4
could be interpreted as having Teff = 60 000 K if it were
assumed to be pure H. This effect is even stronger at higher
temperatures, giving a discrepancy as high as 25 per cent
around 80 000 K. If these levels of He were really present in
the hot DAs, the net effect would be a non-linear compression of the temperature scale towards the hot end. Nevertheless, comparisons between individual objects would still
be valid unless the abundances of helium varied considerably at a given T eff . However, conversely, Napiwotzki (1995)
has demonstrated that the effect trace helium has on the
Balmer lines in the work of Bergeron et al. (1994) may be
exaggerated due to the assumption of LTE. Also, observations of hot white dwarfs with EUVE (e.g. Barstow, Holberg
& Koester 1994a,b, 1995) have shown evidence for a significant He abundance in only one DA, the exceptionally massive star GD50 (Vennes, Bowyer & Dupuis 1996). This
suggests that helium has a minimal influence on the hot DA
temperature measurements and that, on this evidence, pure
H models should be appropriate for determining the DA
temperature scale.
On the other hand, both EUVE and IUE spectra (e.g.
Vennes et al. 1992; Holberg et al. 1993; Dupuis et al. 1995;
Barstowet al. 1996) do show evidence for significant abundances of elements heavier than helium in a number of the
very hot DAs. These elements may well have a similar effect
to helium on the Balmer lines. This is being investigated, but
the evidence is conflicting. Driezler & Werner (1993) found
that with non-LTE model atmospheres, the inclusion of line
blanketing from Fe-group elements had a very small effect
on the hydrogen line profiles. Alternatively, Lanz & Hubeny
(1995) found that the inclusion of heavy elements in their
non-LTE models did have a significant effect on the profiles. Koester (1996) finds that heavy-element blanketing in
LTE models could lower the measured effective temperature of the hot DA G191-B2B by ~5000 K, while a recent
non-LTE study of this star, including an increased level of
Fe line blanketing, shows that the possible effect on the
accuracy of temperature determination could be as large as
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
REJ1847 - 22
REJ1943 + 50
REJ2004 - 56
REJ2009 - 60
REJ2013 + 40%"
REJ2018 - 57
REJ2024 - 42
REJ2024 + 20
REJ2029 + 39
REJ2112 + 50
REJ2127 - 22
REJ2154 - 30
REJ2156 -41
REJ2156 - 54
REJ2207 +25
REJ2210 - 30
REJ2214 -49
REJ2244 - 32
REJ2312 + 10
REJ2324 - 54
REJ2334 - 47
REJ2353 - 24
Alternative
Name
1997MNRAS.286..369M
374 M. C. Marsh et al.
20 , - - - - - , - - - - - - , - - - - - , - - - - - - , - - - - - , - - - - - - , - - - - - - , - - - - - ,
18
REJ0029
16
--.....---:..-.-_~_
14
REJO 138
REJ0623
REJll00
12
" . - - - - - - - _ REJ1148
><
;:J
",------_-lREJ 1257
Q)
:>
......
......,
ro
~---------- REJ/1~3~___
10
Q)
0::
8
~-""""-_ _-1REJ1417
-!REJ1638
,r-'~-_ _ _
6
REJ1726
EJ1800
, . - - - -_ _ _ REJ1847
4
REJ2024.
2
REJ2112
o L -____
3600
~
____- L_ _ _ _
3800
4000
~
_ _ _ _ _ _L -_ _ _ _
4200
4400
~
____
4600
~
______
4800
~
5000
__
~
5200
Wavelength
Figure 1. Representative optical spectra, from both northern and southern hemisphere progranlffies, for the sample of white dwarfs detected
by the ROSAT WFC during the All-Sky Survey.
about 10 per cent (Lanz et al. 1996). In this analysis, we only
use pure H models and, consequently, temperature values
for stars which contain significant quantities of heavy elements may have uncertainties which are much larger than
the formal errors. This potential problem will only apply to
stars with effective temperatures in excess of 55 000 K
We have used the homogeneous and stratified grids of
H + He model atmospheres provided by Detlev Koester
(e.g. Koester 1991; Finley et al. 1997). These span the temperature range 20000-100000 K and logg from 7.0 to 9.0.
Except for those stars known to be DAO white dwarfs, the
analyses were carried out in the limit of negligible helium
abundance or large hydrogen layer mass (MH ), equivalent to
pure H atmospheres. For most of the stars, H p, H y, H~ and
H.lines were included in the analysis. However, several of
the observations obtained by Kidder (1991) spanned only
Hy and H~ lines. Our spectral fitting technique is described
in several earlier papers (e.g. Barstow et al. 1994c), but, as
the results presented in this paper rely on it completely, we
repeat the details here. The analysis was performed using
the program XSPEC (Shafer et al. 1991), which adopts a X2
minimization technique to determine the model spectrum
giving the best fit to the data. Models were convolved with
the appropriate instrument resolution, and all the lines
included were fitted simultaneously, with an independent
normalization constant applied to each, ensuring that the
result was independent of the global slope of the continuum
and reducing the effect of any systematic errors in the flux
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
~
1997MNRAS.286..369M
A UV-selected sample of DA white dwarfs
calibration of the spectrum. The SAAO RPCS, used to
record the southern hemisphere spectra, is a photon-counting instrument. Consequently, the data errors can be
determined simply from the Poisson statistics associated
with the source and background counts in each spectral bin.
However, in dealing with the CCD spectra obtained at the
Steward Observatory, it is difficult to determine errors on
individual points, and we take the following approach. First,
an initial fit to the data is performed without any errors
included. From the scatter on the residuals between the
best-fitting model we estimate the average errors on the
data points, which are included in a subsequent fit. In addition, for both RPCS and CCD data, we correct for any
systematic deviation in the spectral slope between model
375
and data, arising from errors in the flux calibration in the
locality of each line. With the CCD data, a second spectral
fit results in no change to the best-fitting parameters but,
with the inclusion of errors on the data points, yields a lower
reduced X2 • To make sensible estimates of the uncertainty in
the fitted parameters, the value of the reduced X2 should be
less than ~ 2, and this is usually achieved with a single
iteration of the steps outlined.
Two examples of the procedure are shown in Fig. 2
including two contrasting objects, the very hot DA
REJ0029 - 63, observed at SAAO, and the cooler star
REJ0428 + 16 from the Steward sample. Errors on Teff and
logg were determined by allowing the model parameters to
vary until the (jX2 reached the value corresponding to the 10"
(a)
I
I
CJ
""
...".::
~
x
til
V
::s0
CJ
.
~
'0
I
" '"~
y
...0Ei
z
0
'"I
~
'iii"
::s
'0
'!il
10
0
10
I
"...
- --\j~~ !~~~----\I~J~---------------l~b"~~-----
3600
j~I
II I
4000
4200
4400
4600
Wavelength (1\.)
~ jI ~ I I II
4600
5000
(b)
i
«
i CJ
'"I0
....x
'"
"" '"
..."
IV
I
.::
::s
0
0
....
x
.,
~
til
CJ
'0
" '"
.!:!
'iii
Ei...
I
0
....
0
Z
0
'"I
~
'iii"
::s
'0
'!il
10
0
10
I
"...
3600
4000
4200
4400
4600
Wavelength (1\.)
4600
5000
Figure 2. Example spectral fits for two stars where all Hp-H.lines are available. (a) REJ0148 - 25: Teff=24 540 K, logg=7.84 (x; = 1.68); (b)
REJ0029-63: Teff =60595K, logg=7.97 (X;=1.74).
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
V
i
:::: '"
1997MNRAS.286..369M
376 M C. Marsh et al.
Table 2. Summary of derived stellar parameters.
Target
16.82
13.85
15.314
13.37
14.76
13.96
15.87
14.69
14.41
14.79
12.56
14.92
14.83
14.04
13.65
13.951
11.73
17.257
15.815
14.773
16.371
13.032
14.369
15.825
12.089
15.537
8.35
14,82
141178
14.87
14,60
14:22
14.65
14[475
13,190
16 154
14.21
14,93
16.67
16.13
14.455
13.012
14~34
161234
14,64
13:78
14:68
15.773
14.127
15.73
14.33
14.38
13.38
12.99
16.941
15.277
15.083
14.013
11.00
14.83
15.831
15.45
14.606
16.60
14.74
13.949
12.95
13.720
err.
T'II
(K)
I".
bounds
log g
(em s-2)
0.30
0.01
0.011
0.02
0.30
0.03
0.03
0.30
0.02
0.30
0.05
0.30
0.3
0.02
0.05
0.009
0.01
0.013
0.028
0.024
0.012
0.009
0.006
0.017
0.001
0.015
0.1
0.03
0.015
0.04
0.30
0.03
0.30
0.D28
0.020
0.2
0.05
0.30
0.30
0.03
0.020
0.009
0.05
0.D28
0.05
0.30
0.2
0.028
0.028
0.30
0.30
0.05
0.02
0.05
0.020
0.016
·0.028
0.028
0.01
0.05
0.D28
0.05
0.028
0.30
0.04
0.028
0.05
0.D28
46205
47936
60595
34684
28580
44850
38964
24540
30120
63400
62947
31290
32860
42373
34200
58080
57340
53960
31770
45748
51870
32008
70275
33040
62280
41765
24700
32280
44300
53630
37120
51934
29330
36605
23310
33459
53827
31340
36610
34224
44980
22790
28766
47560
29016
27970
39555
38970
55640
31280
25107
55570
37880
49000
42970
34419
46230
38500
24760
36056
37850
54550
880lO
5lO50
43701
45330
28744
31920
44476 - 47850
46839 - 49086
58130 - 63800
34382 - 34998
28320 - 28835
42786 - 48296
38458 - 40708
24300 - 24765
29874 - 30338
62680 - 67280
61481 - 64336
30970 - 31710
32445 - 33135
41577 - 43473
33600 - 34800
55875 - 60170
56040 - 58670
51400 - 58435
31090 - 32160
44360 - 47405
48750 - 56400
31962 - 32287
66400 - 78780
32260 - 34370
60590 - 64140
40330 - 44420
23700 - 25700
31740 - 32840
43310 - 45445
52400 - 54525
36390 -. 37860
48457 - 56007
28970 - 29590
36180 - 37015
22560 - 24120
32882 - 34239
52629 - 55931
30895 - 31680
35780 - 38360
33889 - 34740
44210 - 45310
22570 - 23030
28349 - 29236
46550 - 48630
28750 - 29304
27680 - 28200
38747 - 40503
38620 - 39780
54770 - 56940
30840 - 31770
24764 - 25489
54540 - 56720
37209 - 38584
47000 - 51000
42030 - 44240
33582 - 35135
45540 - 47000
38160 - 38890
24333 - 24805
35479 - 36662
36800 - 38430
49616 - 58835
85400 - 90400
48410 - 53930
42700 - 44759
44600 - 46290
28485 - 29017
31540 - 32260
8.85
7.77
7.97
7.89
7.90
7.96
9.00
7.84
7.70
7.43
7.53
8.44
7.66
9.00
8.80
7.90
7.48
7.62
8.70
7.79
7.29
7.70
7.37
7.80
7.22
8.51
8.65
8.34
7.69
7.64
7.75
8.00
7.79
7.69
7.74
7.86
8.08
7.79
8.69
7.85
7.68
7.57
7.92
7.62
7.79
7.88
7.66
7.91
7.62
8.11
7.81
7.57
7.69
7.70
7.68
7.66
7.70
7.85
7.92
7.71
7.95
8.49
7.79
8.84
7.80
7.73
7.72
8.00
I".
I".
bounds
8.73 7.67 7.76 7.82 7.83 7.70 8.92 7.79 7.62 7.25 7.45 8.35 7.62 8.90 8.52 7.78 7.34 7.33 8.61 7.66 7.02 7.63 7.13 7.57 7.13 8.35 8.35 8.23 7.62 7.54 7.65 7.73 7.75 7.64 7.64 7.70 7.98 7.70 8.50 7.76 7.64 7.53 7.79 7.52 7.71 7.81 7.54 7.81 7.54 8.00 7.73 7.49 7.59 7.50 7.57 7.52 7.62 7.79 7.88 7.61 7.81 8.25 7.70 8.68 7.68 7.68 7.65 7.91 -
8.97
7.86
8.16
7.96
7.97
8.32
9.10
7.92
7.77
7.67
7.62
8.53
7.78
9.10
9.08
8.07
7.58
7.86
8.82
7.90
7.59
7.13
7.50
8.26
7.40
8.65
8.95
8.45
7.80
7.74
7.85
8.22
7.89
7.76
7.83
7.99
8.17
7.89
8.80
7.93
7.77
7.60
8.06
7.70
7.87
7.95
7.78
7.96
7.69
8.21
7.87
7.64
7.79
7.90
7.84
7.82
7.75
7.90
8.09
7.80
8.12
8.77
7.93
9.lO
7.92
7.81
7.79
8,09
Mass(err.)
(Me)
1.16 (0.07)
0.58 (0.04)
0.70 (0.10)
0.61 (0.03)
0.59 (0.03)
0.66 (0.17)
1.22 (0.05)
0.55 (0.02)
0.51 (0.03)
0.52 (0.07)
0.54 (0.03)
0.91 (0.05)
0.50 (0.05)
1.23 (0.05)
1.12 (0.04)
0.66 (0.07)
0.51 (0.04)
0.55 (0.10)
1.07 (0.06)
0.59 (0.05)
0.44 (0.07)
0.51 (0.03)
0.52 (0.07)
0.56 (0.20)
0.46 (0.05)
0.97 (0.07)
1.03 (0.20)
0.85 (0.07)
0.54 (0.04)
0.55 (0.03)
0.55 (0.04)
0.70 (0.13)
0.54 (0.05)
0.52 (0.03)
0.50 (0.04)
0.59 (0.08)
0.65 (0.05)
0.55 (0.04)
1.06 (0.10)
0.58 (0.04)
0.54 (0.03)
0.43 (0.01)
0.61 (0.07)
0.53 (0.03)
0.54 (0.04)
0.58 (0.04)
0.52 (0.05)
0.63 (0.05)
0.55 (0.03)
0.71 (0.06)
0.54 (0.04)
0.53 (0.03)
0.53 (0.04)
0.56 (0.07)
0.54 (0.06)
0.50 (0.06)
0.55 (0.03)
0.60 (0.03)
0.60 (0.09)
0.53 (0.04)
0.64 (0.09)
0.97 (M6)
0.65 (0.21)
1.16 (0.12)
0.59 (0.05)
0.56 (0.03)
0.51 (0.03)
0.65 (0.05)
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
REJ0003 + 43REJ0007 +33
REJ0029 -63
REJ0053 - 32
REJ0108 - 35'
REJ0134 -16
REJ0138 +25
REJ0148 - 25'
REJ0151 + 67
REJ0230 - 47REJ0235 + 03
REJ0237 - 12'
REJ0322 -53
REJ0348 -00
REJ0350 + 17
REJ0457 -28
REJ0505 + 52
REJ0512 -41
REJ0521-10
REJ0550 + 00
REJ0550 -24
REJ0552 + 15
REJ0558 - 37
REJ0605 -48
REJ0623 - 37
REJ0632 - 05
REJ0645 -16
REJ0654 - 02
REJ0715 -70
REJ0720 - 31
REJ0723 - 27"
REJ0827 + 28
REJ0831 - 53REJ0841 + 03
REJ0902 - 04
REJ0907 + 50
REJ1016 - 05"
REJ1019 - 14'
REJ 1024 - 30"
REJI029 +45
REJlO32 +53
REJ1033 -11
REJlO36 +46
REJI043 +49
REJlO44 + 57
REJ1058 - 38"
REJ1100 + 71
REJ1112 +24
REJ1126 + 18
REJ1128 - 02REJ1148 + 18"
REJ1236 +47
REJ1257 +22
REJ1316 + 29!
REJ1340 + 60
REJ1431 + 37
REJ1529 +48
REJ1614 - 08
REJ1623 - 39
REJ1638 + 35
REJ1650 +40
REJ1726 + 58
REJI738 + 66
REJ1746 - 70'
REJ1800 +68
REJ1820 +58
REJ1847 + 01
REJ1847 - 22
mv
1997MNRAS.286..369M
A UV-selected sample of DA white dwarfs
377
Table 2 - continued
Target
14.62
15.05
13.59
14.6
13.61
14.74
16.59
13.38
13.08
14.66
14.17
15.38
14.44
14.58
14.79
11.708
15.66
13.09
15.197
13.441
15.444
I".
err.
10bounds
TefJ
(K)
0.30
33500
0.30
44456
0.30
44200
0.60
47057
0.30
26579
0.30
28597
0.30
50564
0.01
22153
0.02
38866
0.30
48297
0.30
28741
0.30
49764
0.30
45860
24610
0.30
28268
0.30
0.007 65600
31692
0.30
0.01
57380
0.017 45860
0.007 56682
0.014 29120
33330
43381
43650
45680
26034
28230
49332
21749
38139
46455
28363
46513
44600
24490
27823
63790
31312
56190
44490
54883
28720
- 33690
- 45327
C" 44800
- 48230
- 26892
- 28988
- 51810
- 22560
- 39634
- 49934
- 29152
- 53131
- 47200
- 24690
- 28640
- 68510
- 32188
- 58670
- 47320
- 59500
- 29620
log g
(cms- 2 )
7.86
7.54
8.14
7.74
7.78
8.54
7.96
7.79
7.84
7.69
8.18
7.75
7.74
8.16
7.60
7.42
8.07
7.86
7.73
7.64
8.14
I".
bounds
7.82 7.43 8.03 7.49 7.72 8.38 7.83 7.72 7.74 7.56 8.08 7.57 7.64 8.14 7.53 7.24 7.96 7.78 7.61 7.57 8.01 -
7.90
7.66
8.22
8.04
7.85
8.67
8.06
7.87
7.94
7.82
8.26
7.92
7.83
8.20
7.72
7.57
8.17
7.93
7.82
7.78
8.21
Mass(err)
(Me)
0.59 (0.02)
0.49 (0.04)
0.76 (0.06)
0.57 (0.12)
0.53 (0.03)
0.97 (0.10)
0.67 (0.06)
0.52 (0.04)
0.59 (0.05)
0.55 (0.05)
0.75 (0.06)
0.58 (0.07)
0.57 (0.04)
0.73 (0.02)
0.46 (0.04)
0.52 (0.05)
0.69 (0.06)
0.64 (0.03)
0.56 (0.05)
0.56 (0.05)
0.73 (0.08)
Notes.
*mv determined from spectral flux.
·Vmagnitude is taken from Tweedy et al. 1993.
bValues for HZ43 are from Napiwotzki et al. 1993.
CMCf survey (McCook & Sion 1996, in preparation) give V = 14.87, B - V = - 0.18,
U-B= -0.97
dMCf survey (McCook & Sion 1996, in preparation) give V = 15.89, B - V = - 0.30,
U-B= -1.14.
level (bX 2 = 1.0). The errors quoted include only statistical
uncertainties and do not take into account any possible
systematic effects related to the model or observed
spectra.
The values obtained for Teft and logg from the profile
fitting for the whole sample of 89 objects are given in Table
2. The 1CT errors are shown as lower and upper bounds to
each parameter. Also included her~ are the adopted values
for mv for each star in the sample (see Section 4). Where
several photometric observations were available, including
our results and published photometry (available for 47
stars), a weighted average was taken. Objects for which mv
was estimated from the optical spectrum are indicated.
4 OPTICAL PHOTOMETRY
The simple H + He model atmospheres used to interpret
X-ray and EUV data are specified by five main parameters.
These are the effective temperature Teft , the surface gravity,
the interstellar H I column (NH ), the atmospheric helium
abundance (or the mass of the hydrogen layer M H , for a
stratified model) and a distance/radius related normalization constant. Since only three independent data points are
available from the ROSAT survey, additional information is
needed to be able to model the ROSAT data. The optical V
magnitude (mv) may be used as a convenient normalization
constant, while Teft and logg can be derived from the optical
spectra as described above. Magnitude measurements in
other bands are equally useful, although it is mv that is most
often available for catalogued white dwarfs. To provide
optical photometry for the newly discovered white dwarfs
and those known stars for which such measurements had
not been published, an extensive programme of observations has been undertaken to obtain coverage in the U, B
and Vbands. Since the ROSAT survey covers the entire sky,
these observations required time on telescopes in both
heInispheres. Data were obtained for 18 stars in the
southern heInisphere and 25 stars in the north.
The northern heInisphere stars were observed over a twoweek period from 1994 March 28 to April 9 with the 1.0-m
Jacobus Kapteyn telescope (JKT) using an EEV7 CCD chip
placed at the Cassegrain focus, and a set of Harris UBVRI
filters. To increase readout speed, the chip was windowed to
400 x 400 pixel as opposed to the full 1200 x 1200 pixel.
Standards were adopted from the Kitt Peak (Landolt 1983)
and RGO (Argyle et al. 1988) surveys which covered a range
of airmass, 1.0 <X < 2.2. In the main, these standards are
redder than the white dwarf programme stars and while
some blue standards could be observed, giving a complete range in colour of - 0.336 < (B - V) < 1.1 and
"':"'1.245 < (U - B) < 1.1, the observations were not made
under completely photometric conditions. As a result, a
small systematic error was introduced into the determination ofmv.
All the CCD data were reduced using standard techniques, after debiasing· and flat-fielding each frame.
Standard stars at a range of airmass were observed throughout each night and used to determine atmospheric extinction. Since.the character of the atmosphere is not uniform
and, at a given instant, the extinction in various parts of the
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
REJ1943 + 50'
REJ2004 - 56'
REJ2009 - 60'
REJ2013 +40
REJ2018 - 57'
REJ2024 - 42*'
REJ2024 + 20'
REJ2029 +39
REJ2112 + 50
REJ2127 - 22'
REJ2154 - 30'
REJ2156 - 41'd
REJ2156 - 54'
REJ2207 + 25'
REJ2210 - 30'
REJ2214 - 49
REJ2244 - 32'
REJ2312 + 10
REJ2324 - 54
REJ2334 -47
REJ2353 - 24
mv
1997MNRAS.286..369M
378 M C. Marsh et al.
sky may differ, all the standard observations from different
nights were grouped together to give a mean extinction
curve for the whole observing run. Non-uniformity of the
extinction leads to a scatter in the residuals between the
data points and extinction curve. This scatter can be used as
a measure of the internal uncertainty for the magnitudes
( ± 0.03 mag), and dominates the statistical errors. The
derived stellar magnitudes and colour indices were transformed to the standard Johnson photometric system using
the equations of Hardie (1962).
In the southern hemisphere, photoelectric photometry
was obtained as part of a continuing general photometry
programme during 1993-1994 using the VeT photometer,
with GaAs photomultiplier, on the SAAO 0.75-m telescope,
based at Sutherland, South Africa. The data were reduced
using similar techniques (but dealing with photomultiplier
rather than CCD data), using observations of E-region
standards to transform the observations to the standard
Kron-Cousins system (Menzies et al. 1989). The Cousins
system gives magnitudes which are equivalent to those
obtained using that of Johnson for UBVbands.
The combined results for both northern and southern
programmes are listed in Table 3, with the respective observatories indicated. As a consistency check for the JKT pro-
Table 3. UBV photometry of ROSAT DAs.
V
15.314
16.444
13.912
13.827
17.257
15.867
16.198
14.768
16.371
13.043
14.369
15.825
12.089
15.537
15.690
14.178
14.475
13.190
15.747
14.306
14.455
13.028
16.955
16.234
16.781
15.773
14.127
16.809
16.941
15.934
15.272
15.216
15.779
15.083
14.013
15.831
14.606
13.949
13.720
11.708
15.197
13.441
15.444
loB-V
0.011 -0.352
O.OlD -0.307
0.009 -0.359
0.007 -0.258
0.013 -0.266
0.015 -0.215
0.017 -0.296
0.028 -0.298
0.012 -0.351
0.028 -0.239
0.006 -0.331
0.017 -0.214
0.016 -0.329
0.015 -0.262
0.016 -0.306
0.015 -0.351
0.028 -0.305
0.020 -0.168
0.020 -0.265
0.028 -0.033
0.020 -0.276
0.028 -0.181
0.028 -0.333
0.028 -0.504b
0.028 -0.470 b
0.028 -0.291
0.028 -0.280
0.028 -0.396
0.020 -0.289
0.028 -0.379
0.028 -0.305
0.028 -0.151
0.020 0.695
0.028 -1.476 b
0.028 -0.295
0.028 -0.251
0.028 -0.343
0.028 -0.279
0.028 -0.212
0.007 -0.341
0.017 -0.344
0.007 -0.296
0.014 -0.066
U-B
100.015 -1.148
0.021 -1.249
0.011 -1.233
0.019 -1.024
0.017 -0.966
0.015 -1.183
0.023 -1.036
0.022 -1.086
0.017 -1.059
0.022 -1.149
0.003 -1.194
0.024 -1.177
0.021 -1.204
0.020 -1.225
0.021 -1.246
0.020 -0.230
0.022 -1.175
0.016 -0.951
0.016 -1.192
0.022 -1.554 b
0.016 -1.195
0.022 -0.983
0.022 -1.189
0.022 -1.130
0.022 -1.018
0.022 -1.145
0.022 -1.187
0.022 -1.227
0.016 -1.040
0.022 -1.256
0.022 -1.117
0.022 -1.101
0.016 0.206
0.022 -0.863
0.022 -1.141
0.022 -1.269
0.022 -1.250
0:022 -1.194
0.022 -1.123
0.008 -1.205
0.021 -1.161
0.010 -1.168
0.022 -1.175
100.014
0.015
0.009
0.010
0.021
0.013
0.021
0.035
0.015
0.035
0.002
0.023
0.018
0.017
0.018
0.029
0.035
0.025
0.025
0.035
0.025
0.035
0.035
0.035
0.035
0.035
0.035
0.035
0.025
0.035
0.035
0.035
0.025
0.035
0.0315
0.035
0.035
0.035
0.035
0.006
0.020
0.009
0.023
Obs.
SAAO
SAAO
SAAO
SAAO
SAAO
JKT,SAAO
SAAO
JKT
SAAO
JKT
SAAO
SAAO
SAAO
SAAO
SAAO
SAAO (bin)
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT
JKT (bin)
JKT
JKT
JKT
JKT
JKT
JKT
SAAO
SAAO
SAAO
SAAO
Notes
"WD not included in the full sample of 89 objects.
bUnreasonably large value, probable error on magnitude measurement.
(bin) Red colours, possible binary.
© 1997 RAS, MNRAS 286, 369-383
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Target
REJ0029 - 63
REJ0415 - 40 a
REJ0503 - 28 a
REJ0512 - ooa
REJ0512 - 41
REJ0521 -10
REJ0534 - 02 a
REJ0550 + 00
REJ0550 - 24
REJ0552 + 15
REJ0558 - 37
REJ0605 - 48
REJ0623 - 37
REJ0632 - 05
REJ0632 - 57 a
REJ0715 -70
REJ0841 + 03
REJ0902 - 04
REJ0957 + 85 a
REJ1016 - 05
REJ1032 + 53
REJ1033 -11
REJ lD43 + 44 a
REJ1043 +49
REJ1059 + 51 a
REJ1112 + 24 a
REJ1126 + 18
REJ1128 + 17 a
REJ1340 + 60
REJ1425 + 53REJ1431 + 37
REJl440 + 75REJ1501 + 30 a
REJ1529 +48
REJ1614 - 08
REJ1650 +40
REJ1738 + 66
REJ1820 + 58
REJ1847 - 22
REJ2214 - 49
REJ2324 - 54
REJ2334 - 47
REJ2353 - 24
1997MNRAS.286..369M
A UV-selected sample of DA white dwaifs
gramme a white dwarf standard, GD 71, was observed on
one night. GD 71 has mv= 13.032 (Landolt 1992) and we
obtain a value of mv= 15.053, a difference of 0.02 mag,
within our predicted systematic error. Furthermore, the star
REJ0521 - 10 was observed at both the JKT and at the
SAAO. REJ0521 - 10 has a visual magnitude of
15.815 ± 0.028 and 15.889 ± 0.018 from the JKT and SAAO
work respectively, in agreement to within 20'. The value
given for REJ0521 - 10 in Table 3 is a weighted average.
In the sample of stars observed photometrically, it was
not possible to obtain spectra for 13 objects, which were
generally too faint to observe with the 2.0-m class of instruments at our disposal. Consequently, they are not included
in our overall survey of the properties of the EUY-selected
white dwarfs but are listed in Table 3 for completeness, as
no previously published magnitudes exist.
Not all the stars for which spectroscopic data were available
could be included in the photometric programme. However,
useful flux information can be extracted from the spectra
and used to provide an estimate of mV. although the photometric accuracy of the data will not be as reliable, as fewer
standard stars were used and the quality of sky conditions
allowed not so high. The adopted procedure is outlined
below.
Stellar flux is related to a magnitude measurement by the
standard relation
10g/;.(mJ = - OAmx + log/;.(O),
stars observed at Steward. Interestingly, a sample of 12 stars
from the SAAO showed no similar systematic error. Combining these with the corrected Steward values for V. (Fig.
3b) then enables an estimate of the statistical error from the
scatter around the V. = Vp line, yielding 0' = 0.3. This is the
value of the error assigned to all the magnitudes estimated
from the spectral data and included in Table 2.
5 THE MASS DISTRIBUTION OF THE EUVSELECTED SAMPLE
The sample of DA white dwarfs studied in this work were
selected purely on the basis of detection in the EUY by the
WFC during the ROSAT all-sky survey. A bias is clearly
introduced in that only stars with temperatures in excess of
~25 000 K, above which the EUV emission becomes significant, are included. It is important to investigate any effects
of this selection by comparing the sample with other studies,
(a)
1:1
c
's,
o
Cl
rn
o
.......
Cl
Q)
where/;.(O) is the flux for a magnitude 0.0 star in the waveband considered. Values of this constant can easily be
obtained from standard tables (e.g. Zombeck 1990) for all
bands of the Johnson system. While we adopt mv as the
normalization constant for analysis of EUY and X-ray
fluxes, the optical spectra obtained do not cover the wavelength range of the V band but do include the B band.
Consequently, it is possible to determine the value m B and
then estimate mv using B - V measurements obtained from
those stars for which we do have photometric data. At the
temperatures we are considering, the B - V colour depends
primarily on Teff and depends only weakly on gravity (see
Cheselka et al. 1993). As Teff has already been determined
from the Balmer line profiles for all the stars of interest, it is
straightforward to determine B - V from the sample of
Cheselka et a!., which is larger and, therefore, statistically
more reliable than ours. This procedure was carried out for
25 stars where accurate photometry was not available.
Since photometric observations are available for a significant fraction of the stars for which we have spectral data, a
direct test of the reliability of the mv estimates can be made
by applying the technique to these objects. Fig. 3(a) shows
the comparison for 27 stars observed at the Steward Observatory and indicates that the magnitudes determined from
the spectra ev,.) are systematically fainter than the photometric values (Vp). This effect does not appear to be magnitude dependent. Keeping a gradient of 1, a linear fit to the
data gives a constant shift in the V. axis of 0.15 compared to
the Vp = V. line. This is then an indication of the systematic
error which should be subtracted from all values of V. for
....
...+
p., ....
(1)
~
>
V (Photometric)
(b)
1:1
+'
,
c
's,
o
Cl
rn
o
.bCl
Q)
....
p., ....
~
>
V (Photometric)
Figure 3. (a)V, against Vp for 27 stars observed at the Steward
Observatory, showing a systematic error in V,. (b) Magnitudes
from Steward after a systematic correction of - 0.15 mag has been
applied to v" plus 12 stars from the SAAO. Errors on Vp are
indicated. Errors on V, for the SAAO objects are from detector
counting statistics.
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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4.1 Deducing mv from the spectrum
379
1997MNRAS.286..369M
380 M C. Marsh et al.
6 DISCUSSION
To place the mass and gravity distributions into context it is
necessary to compare them with those obtained from optically selected samples. Only then can we assess the possible
biases that might arise from the fact that our sample is
collected on the basis of EUV emission. There are several
studies of mass distributions of large samples of DA white
dwarfs available in the literature. The majority of these have
been summarized by BSL, who compare them with their
own findings. Most of the earlier work is dominated by
objects below Teff ~ 15 000 K and also relies on photometric
methods for deriving temperature and gravity. BSL show
that for Teff > 15 000 K optical Balmer line fitting provides a
far more accurate way of determining Teff and logg, justifying the method used in subsequent studies, including ours.
Reid (1996) has used new gravitational redshift measurements to obtain the masses of 34 field (non-cluster) white
dwarfs in binary systems. These masses, which are independent of spectroscopic mass determination methods, yield a
mean mass of 0.583 ± 0.078 Mo. This is in very good agreement with a redetermination of the BSL spectroscopic
masses using the thick H layer Wood models by Bergeron,
Liebert & Fulbright (1995), who find a mean of
0.590 ± 0.134 Mo.
The study of BSL remains the most accurate and complete sample with which to compare the EUV-selected
25
~
25
.. "':
,,
,
,,
'-
20
20
15
>'-'
<=
>-
~ 15
'-'
<=
Q)
t':?
cr
,,
::l
U.
--,
~
10
10
5
5
, ,
0.4
0.6
9.5
0.8
Mass (Solar Units)
Figure 4. Distribution of white dwarf mass, determined from the
evolutionary models of Wood (1992), for a1l89 stars in the EUVselected sample and all 129 stars in that of BSL. The BSL data are
scaled by the ratio of the sample sizes for ease of comparison.
log g
Figure 5. Distribution of surface gravity for all 89 stars in the
EUV-selected sample and all 129 stars in that of BSL. The BSL
data are scaled by the ratio of the sample sizes for ease of comparison.
© 1997 RAS, MNRAS 286, 369-383
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which may have other biases. One way to do this is by
comparing mass distributions. Once the temperature and
gravity are known, the mass can be found from a theoretical
mass-radius relation. Koester & SchOnbemer (1986)
demonstrated that, when dealing with hot white dwarfs, the
effects of temperature, which cause departures from the
zero-temperature mass-radius relation (Hamada &
Salpeter 1961), should be taken into account. This point is
reiterated in the work of BSL. Failure to do this will lead to
underestimates of the white dwarf mass, because the stellar
radius is assumed to be lower than the true value. In their
analysis, BSL make use of the evolutionary models of Wood
(1992) for DAs with pure carbon cores surrounded by an
envelope with composition log (MHe/M *) = - 4. Subsequently, this work has been extended to deal with thick
(10- 4 Mo) H layers, and these new models, spanning a mass
range of 0.4-1.1 Mo in steps of 0.1 M o ' have been employed
in our work (Wood 1995). A number of stars appear to lie
outside this range, and their masses were estimated by
extrapolating the models. The complete ROSAT sample of
89 stars has a mean logg of 7.899 and mean mass of
0.644 Mo. Mass and logg distributions are shown in Figs 4
and 5 respectively, with the data of BSL, renormalized by
the ratio of the sample sizes (WFC = 89, BSL = 129) to place
them on the same scale, included for comparison.
1997MNRAS.286..369M
A UV-selected sample of DA white dwarfs
Table 4. Comparison of EUY- and optically selected white dwarf
samples.
This work
BSL
Liebert and Bergeron
< logg >
17
< M/M0 >
17
7.90
7.91
7.88
0.38
0.26
0.34
0.644
0.590
0.609
0.189
0.134
0.147
N stars
89
129
200
given mass, higher temperature stars have greater radii and,
therefore, lower surface gravity. For example, a 0.6-Mo star
will have logg = 8.08 at 40000 K, a typical temperature for
an EUV-selected white dwarf, while at 10000 K (a typical
BSL temperature) logg=8.23. However, the mean surface
gravity is very similar to that found by BSL (see Table 4).
The reason for this is apparent from Fig. 5, as discussed
above, with the excess of high-gravity objects among the
ROSAT stars offsetting the lower gravity peak.
The primary difference between the mass distributions of
BSL and this work occurs in the high-mass tail of the two
distributions. As we have shown, the excess in ROSAT stars
above 0.85 Mo is significant at the 99.7 per cent level.
Liebert & Bergeron (1995) discuss the biases which arise in
magnitude-limited optical samples such as the PG survey.
These. samples are biased against high-mass white dwarfs,
because such· stars have smaller radii and thus sample a
smaller spatial volume than less massive white dwarfs of the
same temperature. If EUV luminosity were mass-independent, then this same bias would also apply in addition to the
obvious EUV sampling bias due to the increase of ISM
opacity with distance. However, in comparing the BSL
sample with the EUV-selected objects, we must remember
that we are not considering the same white dwarf temperature range in each case. The BSL mass distribution may not
represent that of an optically selected sample covering the
same temperature range as the EUV sources. Hence the
apparent excess of high-mass stars seen in the EUV data
can only be claimed to exist in comparison with BSL, and
may not be present in a comparison with an optically
selected group spanning the appropriate temperatures.
Unfortunately, a suitable sample does not yet exist in the
literature.
Both Finley (1995) and Liebert & Bergeron (1995) make
the point that there are large differences between the cooling rates of massive and 'normal' ( ~ 0.6 Mo) while dwarfs in
the similar temperature ranges, which affects the number of
stars we might expect to find in a given sample. Indeed, the
sharp decrease in the cooling rate of massive (> 1 Mo)
white dwarfs in the range 40000-50000 K might increase
the numbers of such stars in this interval, compared to the
'normal' (lower mass) population. Alternatively, the excess
may arise because the high-mass white dwarfs are significantly more EUV-Iuminous than lower mass white dwarfs
of the same temperature. Since the photospheric composition of a white dwarf is determined by the balance between
radiative levitation effects and gravity, lower abundances of
heavy elements and, consequently, lower EUV opacity
should be found in those stars with the highest gravities.
However, since many of the high-mass white dwarfs have
temperatures below ~ 40000 K, where heavy elements are
ahnost certainly unimportant (see Barstow et al. 1993), this
does not seem to be a viable explanation of the high-mass
excess.
Whatever selection effect is operating, the presence of a
number of previously unknown massive white dwarfs in the
EUV sample poses several interesting questions. First, has
ROSAT detected a new population of high-mass objects?
Secondly, could the group of stars above ~ 0.9 Mo represent a secondary mass peak produced by coalesced binary
white dwarfs? Such products of binary evolution may be
similar to the excess of low-mass white dwarfs below the He-
© 1997 RAS, MNRAS 286, 369-383
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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ROSAT objects. While both Finley (1995) and Liebert &
Bergeron (1995) have extended the optical work to higher
temperatures than BSL, comparable to the range explored
by ROSAT, these results are preliminary and are not presented in a form that allows detailed comparison. In particular, the mass determinations of Finley (1995) relied on the
zero H layer mass evolutionary models. Consequently, we
compare our sample in detail with that of BSL, but also note
the summarized results of Liebert & Bergeron (1995).
Table 4 lists the mean surface gravities, mean masses and
associated standard deviations of the EUV-selected, BSL
and Liebert & Bergeron samples. The BSL masses were
redetermined using the thick H layer evolutionary models.
The BSL mass and surface gravity distributions are compared with the EUV-selected data in Figs 4 and 5 respectively. To facilitate this analysis, the number frequencies in
each mass/gravity bin of the BSL distributions have been
scaled to account for the differences in the number of
objects included in each sample (89 as compared with 129
objects). It can be seen that the mass distributions have very
similar well-defined peaks, with more than 50 per cent of
stars lying in the range 0.5-0.6 Mo' The ROSAT distribution
has an interesting but modest enhancement in the number
of stars with masses in the range 0.65-0.75 Mo by a factor 4,
but at larger masses (~1 M o )' there seems to be a large
excess of stars compared to BSL. The statistical significance
of this high-mass group can be examined by considering the
BSL sample to represent a parent distribution of objects
and predicting the number of high-mass stars that should be
seen in the EUV-selected group if these stars were drawn
from the same population. Separating the BSL stars into
high- and low-mass groups at 0.85 Mo (124 and five stars
respectively), the probability that an individual star has high
mass ( > 0.85 Mo) is found to be 0.039. After removing from
the ROSAT sample those stars that are already in the BSL
data, we are left with a total of 83 objects, and we predict
that 3.2 ofthese should have high mass. However, 12 stars in
the group fall into this category, an excess of 8.8 compared
to the number expected. Therefore the significance of the
high-mass excess (excess/prediction) is 2.80', corresponding
to a confidence level of 99.7 per cent. We note that a survey
of 18 southern hemisphere EUV-selected hot white dwarfs
by Vennes et al. (1996) reveals the presence of three highmass objects. These are also included in our sample. However, with a comparatively small sample size, Vennes et al.
were unable to construct a mass distribution or determine
the significance of these detections.
Significant differences between the gravity distributions
are apparent in Fig. 5. That of the ROSAT stars peak at
logg=7.75, lower than the BSL group which has a maximum at logg=7.9. Indeed the main portion of the ROSAT
distribution is offset by 0.15 in logg when compared to BSL.
This can be explained straightfolWardly as a result of the
higher temperature of the stars in the ROSAT sample. For a
381
1997MNRAS.286..369M
382 M. C. Marsh et al.
7 CONCLUSION
In this paper we have presented a detailed optical study of
the white dwarfs detected by the ROSAT WFC during its sky
survey and, therefore, selected on the basis of their EUV
flux. The major use of the information on TefI , logg and
visual brightness will be to support a subsequent analysis of
the EUV and X-ray data. However, the distribution of surface gravities and masses derived from the optical work
provides important information about the sample of white
dwarfs and gives some indication of possible biases. The
most striking result is that ROSAT detects a statistically
significant excess in the number of hot, massive DAs that
might be expected on the basis of the optical studies, such as
that of BSL. However, since the optical and EUV samples
do not cover the same range of white dwarf temperatures,
the high-mass excess may arise from differences in the cooling rates of 'normal' (~0.6 Mo) and massive (> 1.0 Mo)
stars. Consequently, this feature may not be a result of
selection on the basis of EUV flux, and might also be
present in an optically selected sample covering an appropriate temperature range. Interestingly, these high-mass
stars may form a discrete population of merged binary white
dwarf systems, but further studies are needed to establish
whether or not this is so.
ACKNOWLEDGMENTS
This paper is based on observations made with the Steward
Observatory 2.3-m telescope operated by the University of
Arizona, the 1.9-m Radcliffe and 0.75-m telescopes of the
South African Astronomical Observatory, and the 1.0-m
Jacobus Kapteyn Telescope operated on the island of La
Palma by the Royal Greenwich Observatory in the Spanish
Observatorio del Roque de los Muchachos of the Instituto
Astrofisica de Canarias. MCM, MAB, MRB, AJP and AES
acknowledge the support of PPARC, UK. JBH acknowledges NASA grants NAGW5-2269 and NAGW5278. The
data reduction and analysis were carried out with Starlink
and NOAO IRAF software. We thank the referee, Dr D.
Finley, for constructive comments on the initial version of
this paper, and for providing additional information which
aided the revision. Finally, we thank Matt Wood for access
to his evolutionary models.
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