METAL LINES IN DA WHITE DWARFS 1 B. Zuckerman. D. Koester. I. N. Reid. and

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1 The Astrophysical Journal, 596: , 2003 October 10 # The American Astronomical Society. All rights reserved. Printed in U.S.A. E METAL LINES IN DA WHITE DWARFS 1 B. Zuckerman Department of Physics and Astronomy and NASA Astrobiology Institute, UCLA, Los Angeles, CA D. Koester Institut für Theoretische Physik und Astrophysik, Universität Kiel, D Kiel, Germany I. N. Reid Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD and M. Hünsch Institut für Theoretische Physik und Astrophysik, Universität Kiel, D Kiel, Germany Received 2003 February 7; accepted 2003 May 31 ABSTRACT We report Keck telescope HIRES echelle observations of DA white dwarfs in a continuation of an extensive search for metals. These spectra are supplemented with new JHK magnitudes that are used to determine improved atmospheric parameters. Of the DA white dwarfs not in binary or common proper motion systems, about 25% show Ca ii lines. For these, Ca abundances are determined from comparison with theoretical equivalent widths from model atmosphere calculations; in a few cases we also obtain Mg, Fe, Si, and Al abundances. If Ca is not observed, we generally determine very stringent upper limits. We compare the data to predictions of previously published models involving the accretion/diffusion of interstellar matter and of comets. The derived abundances are not obviously compatible with the predictions of either model, which up to now could only be tested with traces of metals in helium-rich white dwarfs. By modifying certain assumptions in the published interstellar accretion model we are able to match the distribution of the elements in the white dwarf atmospheres, but, even so, tests of other expectations from this scenario are less successful. Because comet accretion appears unlikely to be the primary cause of the DAZ phenomenon, the data suggest that no more than about 20% of F-type main-sequence stars are accompanied by Oort-like comet clouds. This represents the first observational estimate of this fraction. A plausible alternative to the accretion of cometary or interstellar matter is disruption and accretion of asteroidal material, a model first suggested in 1990 to explain excess near-infrared emission from the DAZ G An asteroidal debris model to account for the general DAZ phenomenon does not presently disagree with the HIRES data, but neither is there any compelling evidence in support of such a model. The HIRES data indicate that in close red dwarf/white dwarf binaries not known to be cataclysmic variables there is, nonetheless, significant mass transfer, perhaps in the form of a wind flowing off the red dwarf. As a by-product we find from the kinematics of GD 165 a likely age of more than 2 Gyr for its probable brown dwarf companion GD 165B. Subject headings: accretion, accretion disks comets: general stars: abundances white dwarfs On-line material: FITS files 1. INTRODUCTION The atmospheres of white dwarfs (WDs) usually show spectral lines of only one element, either hydrogen (spectral type DA) or helium (including spectral type DB). For many years this has been understood as an effect of gravitational separation in the very high gravitational fields in these objects (Schatzman 1947). Elements heavier than helium are observed in a few objects at high temperatures (above 25,000 K), where some elements can be supported in the atmosphere against gravitational forces by selective radiation pressure (see Chayer, Fontaine, & Wesemael 1995, for a detailed discussion). At temperatures below 20,000 K some 1 This paper is based in part on observations obtained at the Calar Alto Observatory of the Deutsch-Spanisches Astronomisches Zentrum and at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA. The Observatory was made possible by the generous financial suppport of the W. M. Keck Foundation. We have made use of the SIMBAD database at CDS. 477 white dwarfs also show traces of heavy elements, although the timescales for gravitational diffusion are short compared to evolutionary timescales (i.e., the heavy elements cannot be primordial) and radiative levitation is negligible. The most widely accepted explanation for these elements is accretion during the passage of the star through an interstellar region of relatively high density followed by subsequent downward diffusion. The elements should thus be visible only during the passage through the cloud and for a relatively short time after the end of the accretion phase. This so-called two-phase accretion scenario has been explored in great detail in a series of papers by Dupuis and collaborators (Dupuis et al. 1992, 1993a, 1993b). Until recently, elements heavier than helium in white dwarfs were seen almost exclusively in DB WDs. For many years G74-7 was the only example of a DA white dwarf with detectable heavy elements such stars are designated DAZ until the discovery of metals in G29-38 (Koester, Provencal, & Shipman 1997) and G (Holberg, Barstow, & Green 1997). There is no obvious reason why

2 478 ZUCKERMAN ET AL. Vol. 596 accretion if that is the correct explanation should preferentially occur for helium-atmosphere objects, and an open question was whether the observational differences between DB and DA WDs could be explained by a combination of several factors: 1. The outer layers of H and He white dwarfs at low temperatures have extended convection zones, where turbulent motions keep the elements homogeneously mixed. Diffusion occurs at the bottom of this convection zone, and the diffusion timescale in the absence of accretion is determined by the diffusion velocity at this layer and the depth of the convection zone. Since the convective layer is much shallower in a DA than in a DB white dwarf, this diffusion timescale is much shorter for a DA. From the figures in Dupuis et al. (1992) we can infer that at T eff ¼ 10; 000 K the timescale is a few times 10 6 yr in the DB WD, but only on the order of 10 4 yr in a DA WD. 2. In the Dupuis et al. model accretion occurs at a high rate during the passage through a dense cloud, assumed to last 10 6 yr in an extremely schematic model following Wesemael (1979), and at a low rate for yr between encounters. In the steady state situation between accretion and diffusion these rates lead to a range of expected abundances for the metals. These expected ranges are very similar for H- and He-rich envelopes, because they are proportional to the diffusion timescale divided by the mass in the convection zone, and the shorter diffusion timescale of a DA is compensated by the smaller mass in the convection zone. 3. During the accretion phase in a cloud roughly for 10 6 yr we thus expect the same metal abundances in H and He envelope white dwarfs. However, after the end of the accretion phase in a DA the abundance will decrease toward the very low and unobservable value for the steady state at low accretion rates on a very short timescale (a few times 10 4 yr). In a helium envelope the metals will stay at relatively high abundances for typically several diffusion timescales (typically a few times 10 6 yr). 4. The comparison is complicated by the very different visibility of metal lines in H and He atmospheres due to the much larger opacity of hydrogen. Dupuis et al. (1993b, their Figs. 3 and 5) estimate that an accretion rate larger by a factor of 10,000 is needed to produce the same 5 Å Ca ii line equivalent width in a DA as in a DB white dwarf. The DA Ca abundances we find in this paper are very similar to those found in DB white dwarfs; the DA equivalent widths, however, are typically må, compared to 10 Å in the DB (note that the equivalent widths do not scale linearly with abundances, since in helium-atmosphere white dwarfs they are on the saturated part of the curve of growth, and in addition the H and K components usually overlap strongly). These faint lines could only be discovered with the help of the new generation of very large telescopes. One conclusion from this discussion is that in the interstellar accretion model we expect to detect metals in DA WDs only during an encounter with an interstellar cloud, when accretion is ongoing, whereas helium-atmosphere objects with metals could be several times 10 6 yr and several parsec away from a cloud encounter. This is especially true for the hotter DA WDs with metals like G29-38 and G238-44, where the convection zones are very small or absent, and diffusion timescales reduce to a few years or even days. This offers the possibility to study the interstellar matter in the neighborhood of the white dwarf and empirically test the accretion/diffusion scenario. Such a search for a correlation between dense interstellar clouds and white dwarfs with metals was inconclusive for helium-atmosphere white dwarfs (Aannestad et al. 1993). Zuckerman & Reid (1998, hereafter ZR98) therefore initiated an extensive search for DAZ white dwarfs using the HIRES spectrograph at the Keck telescope. We report here on a continuation of the search, extending the database to about 100 DA white dwarfs. 2. OBSERVATIONS 2.1. Keck Spectroscopy All of the white dwarfs in the present sample were observed using the HIRES echelle spectrograph (Vogt et al. 1994) on the Keck I telescope. These spectra are available in FITS format in the electronic edition. Some representative spectra appear here as Figures 1 3. Observations were made on the following UT nights: 1997 July 7, 1998 January 22 23, June 24 25, July 21, December 11 12, 1999 April 19 20, July 2 3, and August Excepting the 1999 July 2 3 run, we used the main instrumental configurations described by ZR98: the blue cross-disperser was combined with a 1>15 slit covering the full wavelength region between 3700 and 5200 Å at a resolution of 34,000. In 1999 July the red cross-disperser was combined with the 1>15 slit covering the region between 4300 and 6700 Å at a resolution of 25,000. Our observational program targeted white dwarfs classified as type DA in the McCook & Sion (1999) catalog, although that classification proved erroneous in some cases, as discussed further below. The stars in our sample are selected to have V magnitudes brighter than 16 (for high signal-to-noise ratios) and, primarily, temperatures below 10,000 K. We also observed various known white dwarf/ red dwarf binaries as well as the previously known and relatively hot DAZ WDs G and G Reduction procedures also follow those outlined in ZR98. The spectra were extracted using software routines developed by T. Barlow, producing bias-subtracted, skysubtracted echelle data. The wavelength calibration was defined using the standard routines in the IRAF package, where calibrating thorium-argon arcs were taken at the start and end of each night. Experience has shown that this procedure gives wavelength calibration accurate to better than 3 km s 1, but we tested the accuracy through observations of M dwarfs with accurate radial velocity measurements by Marcy & Benitz (1989). We have not attempted an error estimate for the radial velocity of each target DA white dwarf observed. However, in the notes that follow Tables 1 and 2, we list velocities from Schultz et al. (1996), Maxted & Marsh (1999), and Maxted, Marsh, & Moran (2000). These authors have carefully considered their velocity errors, because the point of their papers was measurement of velocity changes; by comparison with our listed velocities, one can estimate our typical uncertainties. The overall residuals are 0.9 km s 1. Table 1 summarizes results for white dwarfs not known to be in a binary system. Table 2 contains white dwarfs in known or suspected binary systems. The Ca ii K-line

3 No. 1, 2003 METAL LINES IN DA WHITE DWARFS 479 Fig. 1. Representative Ca ii K-line spectra from our HIRES observations. The figure shows three DAZs with intrinsic calcium absorption and one DA (WD ) that displays interstellar calcium absorption. The abscissa is wavelength in angstroms. The spectra are uncorrected for solar motion and are not flux calibrated. The three long horizontal lines along the left-hand side of the figure indicate the zero flux level for the three offset spectra. FITS files for Figs. 1 3 are available in the electronic edition of the Astrophysical Journal. wavelength appears in two orders of the echelle. This line is detected in both orders in all stars we identify as DAZ Infrared Photometry Infrared observations in the J, H, and K passbands were carried out during observation runs in 1999 November (five nights) and 2000 March (three nights) at Calar Alto observatory. Two nights suffered from strong cirrus clouds and were not used for infrared photometry. We used the 2.2 m telescope and the Black MAGIC camera equipped with a Rockwell pixel NICMOS3 HgCdTe detector, yielding a plate scale of 0>64 pixel 1. The moving sky technique was applied by combining five exposures of the same target in order to obtain the sky background. Exposure times typically range from 30 to 75 s for each exposure. Observations of UKIRT faint IR standard stars (Casali & Hawarden 1992) were frequently taken during each night. Data reduction was performed using standard IRAF routines for aperture photometry. Extinction and color coefficients were determined from averages of the most stable nights and treated as constants in the transformation equations. Zero points were determined several times during each night from standard-star observations bracketing the exposures of the program stars. The typical uncertainty is 0.02 for the few stars brighter than magnitude 12 in J, H, and K; it increases to 0.03 in J, 0.05 in H, and 0.03 in K for stars between magnitude 13 and 14; and it is about 0.12 in J, 0.20 in H, and 0.10 in K for stars fainter than 15th magnitude. The resulting magnitudes are shown in Table 3 and may be compared with those listed by 2MASS and in Bergeron, Ruiz, & Leggett (1997) and Bergeron, Leggett, & Ruiz (2001). A summary of such a comparison with the Bergeron papers is presented in the notes to Table 3. For those white dwarfs where we have checked 2MASS, we usually find reasonable agreement with Table 3 magnitudes. For and , where we and Bergeron disagree strongly at K band, the 2MASS magnitudes agree better with ours (see notes to Table 3). 3. ANALYSIS The majority of the observed objects are cool hydrogen-rich white dwarfs with effective temperatures below 10,000 K. Bergeron et al. (1997, hereafter BRL)

4 480 ZUCKERMAN ET AL. Vol. 596 Fig. 2. WD in four orders of the echelle. The abscissa is wavelength in angstroms, and the ordinate is counts at the detector. These spectra are uncorrected for solar motion. In the upper left panel the lines are due to Fe i. In the upper right panel the line near 3821 Å is due to Fe i, while those near 3833 and 3839 Å are from Mg i. The lower left panel shows the Ca ii K-line near 3934 Å, as does the lower right panel. In the latter panel, the Ca ii H-line near 3969 Å and a weak Al i line near 3962 Å also appear. have shown that atmospheric parameters for these objects can be determined very efficiently using optical and infrared (JHK) photometry. The photometry, converted to average fluxes over the wavelength bands, covers a significant part of the complete energy distribution and thus determines T eff. The dependence of the flux distribution on the surface gravity on the other hand is small; this parameter therefore cannot be determined from photometry alone and must be fixed at an assumed value, e.g., 8.00, in the fitting procedure. However, many of these cool objects are not too distant, and therefore reasonably accurate parallaxes are available that can be used to constrain log g. In this study we use a standard 2 method based on the Levenberg-Marquard algorithm (Press et al. 1992) to fit observed to theoretical fluxes. Our procedures are very similar to those of BRL. The solid angle from the radius and distance determines the normalization factor between theoretical and observed fluxes; the radius is connected to the surface gravity used in the fitting procedure via the Wood (1995) mass-radius relation. The theoretical magnitudes used in our analysis have been calculated from our own grid of model atmospheres; the input physics and numerical details are described in Finley, Koester, & Basri (1997). As in BRL, the transmission functions for the filters UBVRIJHK were taken from Bessell (1990) and Bessell & Brett (1988). Integrating the model spectra over the filter bandpasses gives the observable magnitude of a standard star of magnitude zero, except for a normalization constant or zero point. These zero points in all filter bands are determined from a similar integration over the absolutely calibrated flux of Vega, obtained from the STScI Archive, which is based on Hayes (1985), the average IUE spectrum for the UV and an ATLAS12 model for the infrared. Contrary to BRL, who assumed all magnitudes for Vega to be zero, we have taken the average observed magnitudes from the SIMBAD database (0.03, 0.03, 0.03, 0.07, 0.10, 0.02, 0.02, 0.02 for UBVRIJHK). This results in the zero points , , , , , , , for UBV- RIJHK, respectively. Although this is slightly different from

5 No. 1, 2003 METAL LINES IN DA WHITE DWARFS 481 Fig. 3. Close M star white dwarf binary in three orders of the echelle. The axes are the same as in Fig. 2. These spectra are uncorrected for solar motion. See Fig. 2 legend for some line identifications. In addition, an Al i line near Å and a Si i line near 3906 Å can be seen. The double emission line near 3973 Å is due to the Ca ii H line and H in the M star. the procedure used by BRL, they have adjusted their theoretical zero points empirically for observed offsets in the colors; in the end we get almost exactly the same results for the atmospheric parameters of stars, if we use their published magnitudes. In some cases, where we could not obtain a parameter determination from published or our own photometry, or where clearly superior determinations from low-resolution spectroscopy were available, we have used these parameters from the literature (Finley et al. 1997; Bragaglia, Renzini, & Bergeron 1995; Homeier et al. 1998; Leggett, Ruiz, & Bergeron 1998; Vauclair et al. 1997; Koester et al. 1997; Bergeron et al. 1989, 1995a, 1995b, 1997, 2001). The same atmospheric models used for the calculation of magnitudes were also used for the calculation of equivalent widths of metal lines, especially the Ca ii H and K resonance lines, which are the lines most often observed in DAZ white dwarfs. The heavy elements were not considered in the equation of state for the atmosphere models, nor in the calculation of the continuous absorption coefficients. Given the extremely small abundances of these elements, this approximation is completely adequate. After the determination of T eff and log g for each object the equivalent widths of metal lines for these parameters and 12 abundance values from log NðMÞ=NðHÞ ¼ 6:00 to were interpolated from the theoretical tables. The observed equivalent width then was converted to an abundance or upper limit using this table. Metal lines such as the Ca ii resonance lines can also be produced by interstellar absorption. The major tool to distinguish photospheric from interstellar absorption is comparison with Balmer line radial velocities. While the pressure shift of hydrogen lines is negligible, this is not necessarily the case for the metal lines broadened by quadratic Stark effect of electrons and van der Waals broadening by neutral hydrogen. We have very crudely estimated the magnitude of this effect for the Ca ii K line. The Kurucz & Bell (1995) line lists give the broadening parameters log 4 and log 6 as 5.52 and According to the simple classical theory (e.g., Unsöld 1968), the shifts are smaller than the broadening constant by factors of 1.16 and 2.8. Taking the electron and neutral hydrogen densities from our model

6 482 ZUCKERMAN ET AL. atmospheres at an Rosseland optical depth of 0.01, this translates into a shift of 0.22 km s 1 at T eff ¼ 10; 000 K and 4.3 km s 1 at T eff ¼ 6000 K. For most of the observed lines the central parts are formed at higher atmospheric layers with lower pressure, and therefore the shifts are expected to be smaller than 1 km s 1 in most cases. This is true only for hydrogen atmospheres because of the much higher neutral pressure in helium atmospheres the effect will be much larger there. Table 1 summarizes results for normal DA white dwarfs in the sample. The reference code gives the method or the source for the atmospheric parameters as explained in the notes to the table. Our fitting routine also determines errors for T eff and log g, based on the variance matrix of the 2 solution. However, these values depend on the errors of the input data, e.g., the photometry, which is often not known. These formal errors also do not include systematic errors from the models and the transformation of magnitudes to fluxes, and they are therefore in most cases unreasonably small. We do not include these errors in the table, but give an approximate estimate based on the comparison of results from different authors, sometimes using very different methods. Our estimate is that typical errors are about K for T eff below 10,000 K, and up to 500 K for higher values. Errors for surface gravities are typically about 0.1, although they may be significantly larger in some cases with small parallaxes and correspondingly large parallax errors. In a few cases the parallax error seems to be much larger than the catalog value, since the resulting surface gravities would be below 7.0. In these cases we have neglected the parallax and assumed a value of 8.0 for log g. Most of the white dwarfs in Table 1 and Table 2 are in the white dwarf catalog by McCook & Sion (1999), where further data and references can be found. Exceptions are two stars from the Edinburgh-Cape Survey (Kilkenny et al. 1997): WD (EC ) and WD (EC ). About one-third of these white dwarfs were already listed in ZR98. We have included them again in our tables, since for this work we have made new determinations of atmospheric parameters, based partly on new photometric observations. We have also determined as far as possible the surface gravities. Table 1 also includes a few magnetic DA. With the exception of WD (GW ), known to have extremely strong fields, these were analyzed in the same way as the nonmagnetic stars, and the abundances are obviously less certain than for normal DA. A few stars turned out to be not (DA) white dwarfs (see notes to the table). After the completion of our calculation, Bergeron et al. (2001, hereafter BLR) published a paper with new observations and analyses of cool white dwarfs with parallaxes. We have 49 objects in Table 1 in common with that study. For 42 objects the results for T eff and log g agree approximately within the mutual 1 error bars, of which 35 are our own determinations and the rest literature values. Four of the remaining objects are WD , WD , WD , and WD , for which BLR publish new optical and IR colors, whereas we previously used only UBV. We have used their solutions now in the table. For WD BLR obtain a very low surface gravity of log g ¼ 7:20. We had found a similar low value, but since this is very likely a double degenerate, we have used log g ¼ 8 for the determination of T eff and Ca abundance. WD is the well-known double-degenerate G =L We use in our table the result for the brighter (0.5 mag near the Ca ii lines) component of the pair (Bergeron et al. 1989), whereas BLR give the result for the composite object. And finally, for WD our value is taken from Bergeron et al. (1995a), because BLR consider the parallax on which their different result is based to be possibly in error. With a few exceptions, which can easily be explained as due to new and improved observations, this comparison demonstrates that the determination of atmospheric parameters from optical+infrared colors and parallax a method pioneered by BRL is a very stable and reliable method for cool DA WDs. Several objects in our sample have common proper motion companions or were known previously or have been discovered here to have close M star companions, or being TABLE 1 Atmospheric Parameters of Observed White Dwarfs WD Name T eff log g EW [Ca/H] M V H V CaII G <30 < K G <10 < B G <20 < K GD <10 < K G K L <20 < K GD Kiso <30 < K PG <30 < K G <30 < B G <10 < K G K GD <15 < K LHS <20 < K LHS <40 < K GD <10 < K PG HK

7 TABLE 1 Continued WD Name T eff log g EW [Ca/H] M V H V CaII LHS K G BR G <40 < BR G <20 < K Kiso Kiso GD K G <10 < K G <20 < K LHS <15 < K Kiso <10 < K LHS <3 < BB GD K G <10 < K G <15 < K EC <39 < K G <12 < K G <8 < K G <15 < K PG <20 < K G <10 < K EC K G <30 < K G <10 < K ESO BR LP <12 < B EC K EC K G K LHS <35 < K G K G <15 < BR G FK G K PG <10 < K EC <6 < K PG K LHS K GD <15 < K GD <15 < HK GD <10 < K G <10 < K PG < B G K G <5 < B G <15 < K G <10 < K G <20 < K G <12 < K G <25 < BR LHS K G B GD <30 < K G <15 < K G K LHS < GD <5 < BW G <10 < K LTT <10 < BB LTT <3 < BB LHS <50 < BR G <9 < K GD <20 < K GD <30 < B G <15 < K G <15 < K

8 TABLE 1 Continued WD Name T eff log g EW [Ca/H] M V H V CaII LHS <20 < K G KP var var LHS <25 < K LHS <15 < BB Note. EW is the equivalent width in må (or upper limit) for the Ca ii K line, the corresponding abundance is the logarithmic abundance by number relative to hydrogen. The column labeled M gives the method used for parameter determinations or literature source as explained in the sources to the table. Measured heliocentric radial velocities of the H and Ca ii K lines are in units of km s 1. Remarks. Remarks on individual objects; radial velocities are given in parentheses (in units of km s 1 ) : broad H core : no obvious core at H; Maxted et al give (27:0 2:3) : many broad absorption features, possibly subdwarf? Maxted et al give (37:6 1:1) Maxted et al give (20:8 1:2) : prototype DAZ star; radial velocity measurements in 1998 January, and 1999 July and August agree within the errors (see x 2.1) : broad H core; Maxted et al give (26:2 1:4) : very weak H core : radial velocity measurements in 1998 January and December agree within the errors (see x 2.1); other metal lines may be present below 4000 Å : broad shallow Balmer absorption, radial velocity measurements for Ca K line in 1998 January and December agree within the errors : classified DC in McCook & Sion 1999, but clearly shows weak H core : classified DA by Wegner, Africano, & Goodrich 1990; H core is broad; He lines clearly present and therefore a DAB : He lines, broad Ca ii H (308 må), Ca i 4227 (219 må), Mg ii 4482 (475 må), broad Mg i b triplet; probably hot subdwarf : radial velocity measurements in 1998 January and 1999 April agree within the errors; H has a broad core and the Ca ii velocity is probably more reliable : known strongly magnetic DA (Schmidt & Smith 1994), no obvious core at H : individual velocity measurements (H, Ca) are (36.8, 36.1) in 1998 January and (45.5, 41.5) in 1998 December; the table values are averages, but the velocity differences may be real : Maxted et al give (58:6 0:7) : broad H core : radial velocity measurements in 1998 December and 1999 April agree within the errors : Ca ii H (21 må); radial velocity measurements in 1998 December and 1999 April and July agree within the errors : Ca ii H (27 må), Ca i 4227 (16 må); radial velocity measurements in 1999 April and July agree within the errors : broad H core : radial velocity measurements in 1998 June and 1999 April and July agree within the errors : no obvious core at H : radial velocity measurements in 1997 July, 1998 January, and 1999 April and July agree within the errors : no obvious core at H : known magnetic DA (Schmidt & Smith 1994); Hcomponent at , at , þ at : individual velocity measurements (H, Ca) are (21.9, 19.0) in 1998 July and (13.8, 15.5) in 1999 April, also H at 17.5 km s 1 in 1999 July; the table values are averages, but the velocity differences may be real : radial velocity measurements in 1998 June and 1999 April and July agree within the errors : magnetic DA (Bergeron et al. 2001); H , , þ : asymmetric core in H : radial velocity measurements in 1997 July and 1999 April, July, and August agree within the errors; however, the Ca K-line velocity is consistently redshifted from H by about 4 km s : Maxted et al give (27:5 1:0) : Maxted et al give (40:0 1:3) : broad shallow Balmer absorption : Maxted & Marsh 1999 give (1:9 0:8) : Maxted & Marsh 1999 give (5:8 1:1) : Maxted & Marsh 1999 give ( 5:9 0:9) : radial velocity measurements in 1999 April, July, and August agree within the errors : GW , well known magnetic object : Maxted & Marsh 1999 give (53:0 0:6) : broad square core, known magnetic DA (Koester et al. 1998) : Maxted & Marsh 1999 give (75:4 0:2) : no obvious core at H : Maxted et al give ( 2:4 0:3) : broad H core : listed as radial velocity variable by Saffer, Livio, & Yungelson : (44.6, 47.0) in 1997 July, (41.7, 35.0) in 1998 December, (36.7, 36.1) in 1999 August and (37.9, 35.2) in second spectrum of 1999 August; H at 36.2 km s 1 in 1999 July; Maxted & Marsh (1999 give (44:7 0:7) in 1997 August and 1998 June; ZZ Ceti star; second DAZ to be detected (Koester et al. 1997) : very strange H core, noted before by Koester et al and Maxted & Marsh 1999, who give (43:6 0:6). Sources. Sources of atmospheric parameters or method used to determine them (Kx) in case of new determinations for this paper: K1: optical magnitudes, parallaxes, and our own JHK observations; K2: optical and IR magnitudes from literature+parallaxes; K3: optical and IR magnitudes, no parallaxes, log g ¼ 8:00 fixed; K4: UBV(RI )+parallax from literature; K5: UBV(RI ), no IR, no parallax, log g ¼ 8:00 fixed; K6: estimate from B V only, log g ¼ 8:00 fixed; K7: T eff estimated from U B; logg ¼ 8:00 fixed; K8: T eff estimated from Strömgren colors; log g ¼ 8:00 fixed. BR: Bergeron et al. 1997; FK: Finley et al. 1997; HK: Homeier et al. 1998; BB: Bragaglia et al. 1995; BW: Bergeron et al. 1995b; KP: Koester et al. 1997; B01: Bergeron et al. 2001; B03: Bergeron 2003, private communication.

9 METAL LINES IN DA WHITE DWARFS 485 double degenerates. Observational data and analysis results for the DA WDs in these systems are collected in Table 2. If the atmospheric parameters for the double degenerates or candidates have been determined in this work (code Kx), the fit to the optical and IR colors as well as the abundance determinations proceeded in the same way as for single DA WDs. For the DA+dM pairs we have in most cases used the U B color, which should be least contaminated by the red companion. However, in all these cases the abundances or limits derived should be used with some caution Other Elements Several objects with Ca ii lines show other metal lines in addition; results for these are collected in Table 4. The abundances of iron and magnesium are typically larger than those of calcium by about a factor of 10 20, with the only exception of WD , where the Fe lines are very weak (5 10 må) and the uncertainty of the abundance is large. In addition to the elements listed in Table 4 which are present in our standard blue cross-disperser setup, the sodium D-line frequencies were included in our 1999 July red cross-disperser setup. Of the dozen white dwarfs observed in 1999 July, interstellar sodium was detected toward WD (Table 5). None of the other white dwarfs showed evidence of the D lines. The significance of this nondetection is discussed in xx 4.3 and The ZZ Ceti Star GD 165 As an aside to our primary project of searching for DAZs among cool white dwarfs, we measured the radial velocity of the white dwarf GD 165 (WD ), which is a warm ZZ Ceti star with an L-type companion GD 165B (Kirkpatrick et al. 1999). Since its discovery as the first L- type object known (Becklin & Zuckerman 1988), the status of GD 165B minimum-mass main-sequence star or highmass brown dwarf has been uncertain. Knowledge of the age of the system will help to place it in the T eff versus age plane along with theoretical models for very low mass stars and cooling brown dwarfs (e.g., Fig. 8 in Kirkpatrick et al. 1999). Among various ways to study GD 165B, one can use the observed space motion to estimate a likely lower limit to the age. On the basis of its H and H lines, we measured a radial velocity for GD 165A equal to 0 km s 1. GD 165A is an average-mass white dwarf with M 0:6 M (Bergeron et al. 1995b), and we adopted a gravitational redshift of 30 km s 1 in a calculation of UVW ( 25, 34, 21). This corresponds to velocities of ( 15, 29, 14) relative to the local standard of rest (LSR), or a total space motion of 35.5 km s 1. In principle, any individual star of any age can have any velocity. However, comparing GD 165B against measurements of the velocity dispersions of stars of known ages (Wielen 1974) shows that the observed space velocity corresponds best with a star of intermediate age, that is, most likely in the range from about 2 to 5 Gyr. This age agrees well with that deduced by Kirkpatrick et al. (1999) from considerations of the mass of GD 165A and corresponding stellar evolutionary timescales Interstellar Calcium and Sodium Because the velocity of most stars in Table 1 differs substantially from the velocity of interstellar gas within 60 pc of the Sun (see, e.g., Table 6), in principle it is possible to detect interstellar calcium and sodium along most white dwarf sightlines. However, because of the proximity to Earth of our sample and the low density of local gas, we did not expect to detect interstellar lines. Thus, we were surprised to see interstellar lines toward GD 31 (WD ), LHS 1549 (WD ), and G140-2 (WD ). Ca ii and Na i seen toward these stars is evidently interstellar because the line radial velocities differ greatly from the H line velocity. Also, in the case of double degenerate LHS 1549, due to orbital motion, the H line velocity is variable while the Ca ii K line has constant velocity (see Table 5). Distances to WD and WD are 17 and 43 pc, respectively. The calcium lines are strong compared to calcium seen along sightlines toward typical nearby TABLE 2 DA White Dwarfs Known or Suspected to Be in Binary Systems WD Name T eff log g EW [Ca/H] M Binary? V H V CaII Obs G K4 DA+dM 20.7 (em) 22.0(em) G <10 < K1 dd G <10 < K3 dd G <10 < BF dd 16.6/ G <5 < K6 cpm G LRB dd LHS <3 < BL dd Dec Aug GD <20 < K9 DA+dM Rubin K5 DA+dM 69.7(em) LHS K4 DA+dM 27.1(em core) HZ DA+dM G <30 < K1 cpm GD <15 < K1 dd 38/ LTT DA+dM G <10 < K1 cpm PG <15 < K7 DA+dM PG K7 DA+dM? Dec Apr Apr

10 TABLE 2 Continued WD Name T eff log g EW [Ca/H] M Binary? V H V CaII Obs PG K7 DA+dM 45.2(em) 9.4? 1998 Dec (em) Apr G <15 < BL cpm PG K7 DA+dM 23.3(em) 66.7? Case SGW DA+dM 141.7(em) LHS <40 < BR cpm EC DA+dM G <7 < K1 dd EC DA+dM (em) EC DA+dM (em) GD <30 < K7 dd? 106/159? LTT <3 < VS cpm Case <10 < K1 dd Wolf <20 < BL cpm G <8 < K4 DA+dM KPD K7 DA+dM 12.0(em) GD K7 DA+dM (em) LP <5 < K1 cpm Note. The meaning of the columns is the same as in Table 1, except for an additional column with the binary status. dd: double degenerate system, DA+dM: red dwarf/white dwarf system, cpm: common proper motion companion. If several velocity measurements are available, the dates of the observations are given in the last column. Radial velocities are in units of km s 1. Remarks. Remarks on individual objects; radial velocities are given in parentheses (in units of km s 1 ); (em) indicates that the H velocity is that of the red dwarf in DA+dM systems, exceptions are listed as (em core) : composite spectrum with emission lines originally noted by Probst & O Connell 1982 and also by Schmidt & Smith Schultz et al. 1996: spatial separation of DA and dm is medium. They give (17:6 0:4) for the velocity. Marginal Ca ii absorption? : velocity variable according to Maxted et al with likely period near 6.4 or 1.2 days; suspected as double degenerate by Bergeron et al from colors and H line profile fit : listed as known or suspected binary by Bergeron et al. 2001; Maxted et al give (7:1 0:4) : This is the well know double degenerate object G or EG 11 (Bergeron et al. 1989). The parameters in the table are those for the brighter component. We leave this as the only certain double degenerate in this table since both components have well determined atmospheric parameters : Ca ii lines double; close binary based on large velocity separation between the two components of the Ca ii line; Ca abundance estimated assuming 1 2 of the EW in the table for each component; classified DZ9 by Hintzen & Strittmatter 1974; H detected subsequently, for example, by Bergeron et al. 1997, who classify G77-50 as DZA. H line clearly present in LRIS spectrum obtained on the Keck II telescope by R. Becker, private communication. In 1998 December we measured 28.8 and 62.8 km s 1 for the radial velocities of the two stars, and in 1999 August we measured 49.4 and 84.0 km s : Ca ii interstellar; suspected double degenerate (Bergeron et al. 2001), confirmed by H variation; see also Table : DA+dM binary (Koester et al. 2001) : known DA+dM binary Rubin 80; Schultz et al. 1996: separation medium : LHS 1660, DA+dM, 7 hr period; emission lines from companion, many metal features. H velocity from emission core in absorption line, M dwarf emission at 255 km s 1 ; accretion disk/hot spot on WD? : well known binary in Hyades, 13 hr period, emission lines : broad H core : double H cores : known close binary (Schultz et al. 1996; Maxted & Marsh 1999). Maxted et al give large (180 km s 1 ) velocity separation between emission and absorption components of H. These agree reasonably with our measurements taken a few months earlier (v em ¼ 155 from H and Ca ii) : separation not close (Schultz et al. 1996); they give the H emission velocity ( 5:8 1:3) : 14 hr period; H partially filled in, Ca ii emission at 49.4 (1998 December); Ca ii emission 39.3, H emission 71.4 (1999 April); Ca ii emission 97.8, H emission 98.8 (1999 April) : Ca ii absorption : classified binary in McCook & Sion 1999; separation not close, velocity ( 19:5 0:6) (Schultz et al. 1996); Ca ii emission at 19.2 km s : Case 1, DA+dM, 16 hr period : no core in H : newly identified DA+dM binary : dd according to Bergeron et al : newly identified binary with Ca i,caii, and Mg i emission lines : newly identified binary : two H components, possible newly identified double degenerate : asymmetric core in H : broad core H and H; likely dd according to Bergeron et al. 2001; this star is in the ZZ Ceti temperature range, where Koester et al found unusual H core profiles : Maxted & Marsh 1999 give ( 15:7 0:7( : Maxted & Marsh 1999 and Maxted et al see emission slightly to the blue of the absorption core. Our data confirm this with a separation of 37 km s 1 ; also discussed in Schultz et al. 1996, with no detectable H emission and a? for binarity; over a 10 month period radial velocity measurements by Maxted & Marsh 1999 and us indicate a velocity constant to a few km s 1, which is strong evidence that this is not a close binary : close binary according to Schultz et al : close binary according to Schultz et al. 1996; He, H, and metal emission; Na i absorption from the white dwarf is less than 5 må,h core partially filled in : Maxted et al give (7:5 1:7). Sources. The codes for the method used to determine atmospheric parameters are the same as for Table 1; the only new ones are K9: estimate from fit to Å of Koester et al. 2001, log g ¼ 8:00 fixed; BF: Bergeron et al. 1989; SGW: Sion, Guinan, & Wesemael 1984; BL: Bergeron et al. 1995a; VS: Vauclair et al. 1997; LRB: Leggett et al

11 TABLE 3 JHK Magnitudes of the Observed White Dwarfs WD J H K : : : : 16.28: 15.30: : : 14.14: : : : : TABLE 3 Continued WD J H K : : : 14.59: : EC : 15.21: 15.37: EC : 12.53: EC : 16.07: 15.58: Note. A colon denotes an uncertainty larger than The table contains 56 white dwarfs, of which about half appear in Bergeron et al or Bergeron et al with infrared photometry. In general, our photometry and their s agrees to 0.1 mag or better. One exception is for which our K magnitude of disagrees badly with the given in Bergeron et al and also with the given in Bergeron et al So, we checked the 2MASS database. Surprisingly, there are two stars, separated by about 40 00, that might be the infrared counterpart of ; the 2MASS K magnitudes of both stars are Another exception is , for which we find K ¼ 13:52 and Bergeron et al give K ¼ 13:79. The 2MASS value is TABLE 4 Abundances of Elements Other than Calcium WD Mg/H Mg/Ca Fe/H Fe/Ca Al/H Al/Ca Si/H Si/Ca Solar : 84: a 21 a b 0.46 b Note. The first column for each element gives the logarithmic number ratio relative to hydrogen, the second is the number ratio of the element relative to Ca. a From Holberg et al. 1997, based on Mg ii b From Holberg et al. 1997, based on Si ii 1309 and Mg abundances are from 3832 and 3838 except for and , for which 4481 was used. Fe abundances from 3820 and Si abundances from Al abundances from 3944 and