aa r X i v : . [ a s t r o - ph . S R ] N ov Hot White Dwarfs
Edward M. Sion
Department of Astronomy & Astrophysics, Villanova University
ABSTRACT
The article covers the physical properties and evolution of single white dwarfsranging in temperature from 20,000K to 200,000 and higher, the hottest knowelectron-degenerate stars. After discussing the classification of their spectra, theauthor reviews the known properties, parameters, evolutionary state, as well aspersisting and new puzzles regarding all spectroscopic subclasses of Hot WhiteDwarfs: the hot DA white dwarfs, the DAO white dwarfs, the PG1159 degen-erates, the DO white dwarfs, the DB white dwarfs, the DBA white dwarfs, andthe Hot DQ white dwarfs (an entirely new class). The most recent observationaland theoretical advances are brought to bear on the topic.
1. Introduction
Research on hot white dwarfs during the past thirty years has greatly expanded, as manynew discoveries, and the new questions they raise, have emerged from increasingly larger,deeper surveys conducted with multi-meter class ground-based telescopes, the
InternationalUltraviolet Explorer ( IUE ), the
Hubble Space Telescope ( HST ), the
Extreme Ultraviolet Ex-plorer ( EUVE ), and the
Far Ultraviolet Spectroscopic Explorer ( FUSE ). This review will fo-cus on white dwarfs ranging in temperature from 20,000 K up to 200,000 K and higher, whichare the hottest white dwarf stars known. Since the mid-20th century, the earliest spectro-scopic surveys of white dwarf candidates from the proper motion selected samples of WillemLuyten and Henry Giclas were carried out by Jesse Greenstein, Olin Eggen, James Liebert,Richard Green, and others. The selection criteria employed in many of these surveys did notreveal a large number of hot white dwarfs because the surveys lacked ultraviolet sensitivityand also missed objects with low flux levels in the optical. Nevertheless, the earliest surveysquickly revealed that white dwarfs divide into two basic composition groups, with hydrogen-rich (the DA stars) and helium-rich atmospheric compositions (the DB and other non-DAstars). The origin of this dichotomy still represents a major unsolved problem in stellarevolution, although theoretical advances in late stellar evolution made starting in the 1980s,as well as advances in modeling envelope physical processes and mass loss, have shed impor-tant new light on this puzzle (Fontaine & Michaud 1979; Vauclair, Vauclair, & Greenstein 2 –1979; Iben et al. 1983; Schoenberner 1983; Iben 1984; Chayer, Fontaine, & Wesemael 1995;Unglaub & Bues 1998, 2000).The spectroscopic properties of white dwarfs are determined by a host of physical pro-cesses which control and/or modify the flow of elements and, hence, surface abundances inhigh gravity atmospheres: convective dredge-up, mixing and dilution, accretion of gas anddust from the interstellar medium and debris disks, gravitational and thermal diffusion, ra-diative forces, mass loss due to wind outflow and episodic mass ejection, late nuclear shellburning and late thermal pulses, rotation, magnetic fields, and possible composition relics ofprior pre-white dwarf evolutionary states. Virtually all of these processes and factors mayoperate in hot white dwarfs, leading to the wide variety of observed spectroscopic phenomenaand spectral evolution.The basic thrust of research on hot white dwarfs is three-fold: (1) to elucidate theevolutionary links between the white dwarfs and their pre-white dwarf progenitors, whetherfrom the asymptotic giant branch (AGB), the extended horizontal branch, stellar mergers, orbinary evolution; (2) to understand the physics of the different envelope processes operatingin hot white dwarfs as they cool; and (3) to disentangle and elucidate the relationshipsbetween the different spectroscopic subclasses and hybrid subclasses of hot white dwarfs asspectral evolution proceeds. This includes the source of photospheric metals, the chemicalspecies observed, and the measured surface abundances in hot degenerates. The evolutionarysignificance of certain observed ion species is complicated by the role of radiative forces andweak winds in levitating and ejecting elements at temperatures >
2. Remarks on the Spectroscopic Classification of Hot White Dwarfs
Before discussing the hot white dwarfs, a brief discussion of their spectroscopic clas-sification is appropriate. The non-DA stars fall into six subclasses, including the PG 1159stars (75,000 K < T eff < < T eff < < T eff < T eff < T eff , from 4000 K up to 120,000 K and higher. Figure 1shows representative optical spectra of several DA white dwarfs. Hot DA stars that containdetectable helium are classified as DAO if He II is present and DAB if He I is present. Becauseof the importance of temperature as a direct luminosity and age indicator in white dwarfs,and the fact that white dwarfs span enormous ranges of T eff (e.g., the H-rich white dwarfsspan a temperature range from 4500 K to 170,000 K!), a temperature index was introducedby Sion et al. (1983) defined as 10 × θ eff (= 50 , /T eff ). Thus, for the hot DA and non-DAstars, their spectral types can be expressed in half-integer steps as a function of temperature;for example, the DA sequence extends from DA.5, DA1, DA1.5, DA2, DA2.5...DA13. A DA2star has a temperature in the range 22,400–28,800 K, while a DA2.5 has T eff in the range18,327–22,400 K. Similarly, the sequence of DB stars extends from DB2, DB2.5, DB3, DB3.5,and so on. A DB2 star has a temperature in the range 22,400–28,800 K, while a DB2.5 has T eff in the range 18,327–22,400 K. Figure 2 shows representative optical spectra of severalDB white dwarfs. For the hot DA stars ( T eff > g = 7, which is traditionally adopted as theminimum defining gravity for classification as a white dwarf star (versus a high gravity sub-dwarf; Greenstein & Sargent 1974); and (2) the assignment of the primary spectral class fora white dwarf is determined by the element represented by the strongest absorption featuresin the optical spectrum. However, by this criterion, the PG 1159 stars, for which eitheroxygen (e.g., O VI) or carbon (e.g., C IV) are the strongest optical lines (with He II featuresweaker), should be classified as DZQO or DQZO depending upon whether O VI or C IV arestrongest, respectively. Since many of these objects have atmospheric compositions whichare not completely helium-dominated (cf., Werner & Heber 1991; Werner et al. 1991), it isinappropriate to assign spectral type DO on the Sion et al. (1983) scheme since the primaryO-symbol denotes a helium-dominated composition. Hence, the primary spectroscopic typeis adopted here as the atom or ion with the strongest absorption features in the optical spectrum, where applicable (for example, it is possible that the strongest absorption featuremay lie in the ultraviolet). This is the scheme used for classifying the hottest degenerates. 4 –In practice, a degenerate classification is withheld for any PG 1159 star with log g <
7. However, these objects are designated PG 1159 as given by Werner & Heber (1991)and Werner et al. (1991), and subsequently used by Napiwotzki & Schoenberner (1991) andDreizler, Werner, & Heber (1995). These designations are: E for emission, lgE for low grav-ity with strong central emission, A for absorption, Ep for emission/peculiar, and EH, lgEHor AH for hybrid PG 1159 stars that have detectable hydrogen. The temperature indexwould differentiate the hot C-He-O stars from the well-known, very much cooler DZ and DQdegenerates below 10,000 K. For example, PG 1159 itself (log g = 7, T eff = 110,000 K, C IVabsorption as the strongest optical lines) would be classified as DQZO.4. The obvious disad-vantage is the inevitable confusion with the cool DQ and DZ degenerates in cases for whichthe temperature index is missing or there are no He II absorption features (e.g., H 1504;Nousek et al. 1986). In a case like H1504, where no helium is detected, T eff = 170,000 K,and log g = 7, the classification DZQ.3 is assigned.The hot DQ stars are an entirely new subclass of hot white dwarfs. Unlike the previouslyknown, and cooler, DQ white dwarfs which have helium-dominated atmospheres, the hot DQstars have atmospheres that are dominated by carbon! These hot DQs probably evolve froma different progenitor channel than the cool DQs. While the temperature index can beused to distinguish these C-dominated objects from the cooler He-dominated DQ stars, theiruniqueness suggests a special classification designation as “hot DQ” stars defined by thedominant presence of C II features in their optical spectra. There is certainly a precedentfor having a special designation for this unique class of C-dominated white dwarfs. ThePG 1159 stars merited their own special designation, rather than classifying them as DZ,DQ, or DO based upon whether O, C, or He features dominated their optical spectra. Thedesignation PG 1159 is widely used to distinguish these unique, exotic objects from helium-dominated DO stars at lower temperatures. See the chapter by P. Dufour in this volume fordetails of the hot DQ white dwarfs.
3. The Hot DA Stars
The total number of hot ( T eff > T eff and estimate log g in the far-ultraviolet with Orfeus (e.g., Dupuis et al. 1998),the Hopkins Ultraviolet Telescope (HUT; e.g., Kruk et al. 1997), and FUSE (Sahnow et al.2000), there are two different widely known discrepancies that plagued the reliable deter-mination of hot DA physical parameters: (1) an inability to fit all of the Balmer lines si-multaneously with consistent atmospheric parameters (the so-called Napiwotzki effect; cf.Gianninas et al. 2010 and references therein); and (2) the disagreement between the parame-ters derived from fitting optical spectra and those derived from fitting far-ultraviolet spectra(e.g., Finley, Koester, & Basri 1997 and references therein). The Napiwotzki effect has beenresolved by adding metals (not detected in the optical spectra) to the model atmospheres,which provides a mild back-warming effect. The fact that the analysis of far-ultravioletspectra from the FUSE archive reveals a correlation between higher metallic abundances andinstances of the Balmer line problem strongly supports this scenario (Gianninas et al. 2010).However, the disagreement between parameters derived from optical and far-ultraviolet spec-tra remains.A large fraction of the hot DA stars observed in the extreme- and far-ultraviolet haverevealed trace abundances of numerous heavy elements which are presumably radiativelylevitated against downward diffusion by radiative forces (Chayer et al. 1995). Extreme-ultraviolet observations of hot DA white dwarfs have been particularly effective in revealinglevitated trace metals in their atmospheres (Finley et al. 1997). This occurs because theextreme-ultraviolet flux of a hot DA star can be strongly suppressed due to both the lowopacity of the residual neutral hydrogen shortward of 300 ˚A, and the strong continuum ab-sorption and heavy line blanketing in that same extreme-ultraviolet wavelength range due toany trace metal ion constituents that may be present in the photosphere. Finley et al. (1997)point out an extensive literature on extreme-ultraviolet analyses of hot DA white dwarfsincluding, for example, work by Kahn et al. (1984), Petre, Shipman, & Canizares (1986),Jordan et al. (1987), Paerels & Heise (1989), Barstow et al. (1993), Finley et al. (1993),Jordan et al. (1994), Vennes et al. (1994), Vennes et al. (1996), Wolff, Jordan, & Koester(1996), and Marsh et al. (1997).Absorption features due to C, N, O, Si, Fe, and Ni, have been seen, and in one object,PG 1342+444, absorption lines of O VI are detected which had previously only been seen inthe subluminous Wolf-Rayet planetary nuclei and the PG 1159 stars (Barstow et al. 2002).On the other hand, there are also a sizable number of hot DA stars that appear to be metaldeficient since radiative forces theory (Chayer et al. 1995) predicts that any metals present 6 –should be levitated.There are now well over 100 known DA white dwarfs with T eff > − T eff estimates in the range 65,000 K – 72,000 K and log g in the range 8.4–9.1, correspondingto a mass of M wd ≈ M ⊙ (Finley et al. 1997; Vennes et al. 1997; Dupuis et al. 2002;Koester et al. 2009); and (ii) the fact that the DA stars reach temperatures above 100,000 K,with the hottest currently known DA star, WD 0948+534, at T eff & ∼ > T eff > ∼ M bol = 2, which they interpret as atransition of non-DA atmospheres into DA atmospheres during white dwarf evolution. Thetransition would occur as trace amounts of hydrogen float to the surface and give rise toH-features in the optical spectrum.It is well known that the hottest DA stars are cooler than the hottest non-DA stars(Sion 1986; Fontaine & Wesemael 1987). Unless some fraction of the PG 1159 degeneratesundergo spectral evolution into hot DA stars when previously “hidden” hydrogen floats up totheir surface, then the progenitors of the hottest DA stars represent a separate evolutionarychannel. This would be contrary to the single channel scenario of Fontaine & Wesemael(1987) In the Hertzsprung-Russell (H-R) diagram, the hottest DA white dwarfs appear toconnect up with the H-rich central stars of planetary nebulae. There is very likely a directevolutionary connection. 7 – The DAO stars are hot DA stars that exhibit ionized helium in their optical spectra.They are characterized by surface temperatures in the range 50,000 K < T eff < g < .
5, and log [He / H] = − not be in binaries and, hence, maybe the descendants of extended horizontal branch progenitors. Because pure DA and DBstars have monoelemental atmospheres, the hybrid composition DAO stars offer insights intowhite dwarf spectral evolution.There is typically a large discrepancy in the T eff values determined for DAO white dwarfsfrom optical Balmer line spectra compared to temperatures derived from Lyman line spectraobtained with FUSE (see Table 4). This behavior echoes the Balmer line versus Lyman linetemperature discrepancy noted by Barstow et al. (2003) for the hot ( T eff > M wd > . ⊙ appearto represent an evolutionary channel connecting them to the AGB stars, since DAO whitedwarfs more massive than 0 . ⊙ would have been massive enough for helium shell burningon the horizontal branch. In this scenario, the DAO stars could also be the progeny of H-richplanetary nebula central stars or even hybrid PG 1159 stars (containing some H) in whicha “hidden” reservoir of hydrogen floats up to the surface. Recently, Gianninas et al. (2010)contend that the post-extreme horizontal branch evolution is no longer needed to explainthe evolution for the majority of the DAO stars, and that the presence of metals might drivea weak stellar wind, which, in turn, could explain the presence of helium in DAO whitedwarfs. Nevertheless, it is still not possible to definitively establish these different potentialevolutionary links.
4. The PG 1159 Stars
The most exciting stellar discovery of the Palomar Green survey was a class of extremelyhot, high luminosity degenerate objects known as the 1159 stars (Green, Schmidt, & Liebert1986). Subsequently, large surveys (Palomar Green; Hamburg-Schmidt; Hamburg ESO,Wisotzki et al. 1996; and, most recently, the SDSS) have uncovered the majority of theknown PG 1159 stars . The PG 1159 stars reveal spectra typically devoid of hydrogen. In-stead, they are dominated by He II and highly ionized, high excitation carbon, especially abroad absorption trough in the region of 4670 ˚A comprising He II λ T eff as high as 100,000 K The most recent and only discovery within the last 10 years, besides those from the SDSS, isHE 1429 − HST in order to acquirespectra of sufficient quality. They carried out state-of-the-art analyses of the hottest (pre-)white dwarfs by means of NLTE model atmospheres, which include the metal-line blanketingof all elements from hydrogen to nickel (Rauch & Werner 2010).The spectral analysis of the PG 1159 stars revealed a range of temperatures of T eff =75,000–200,000 K and gravities of log g = 5.5–8.0 (Rauch et al. 1991; Dreizler et al. 1994;Werner et al. 1996). Prior to the SDSS, only 28 PG 1159 stars were known. From theempirically derived ranges of their parameters, it was obvious that the position of someof the PG 1159 stars (specifically the lower gravity members, subtype lgE; Werner 1992)in the H-R diagram overlapped the hot central stars of planetary nebulae. The remainderof the PG 1159 stars are more compact objects with higher (white dwarf) surface gravities(subtype A or E). Thus, the PG 1159 stars appear to be evolutionary transition objectsbetween the hottest post-AGB and white dwarf phases. Figure 3 shows a summary plot oflog g versus log T eff from the most recent analyses of the PG 1159 and hot DO white dwarfs.For comparison, evolutionary tracks by Althaus et al. (2009) are also plotted.Until very recently, Fe lines had never been detected in PG 1159 stars even though Felines are seen in some hot DA, DO, and DAO white dwarfs. In the far-ultraviolet, featuresdue to Fe VII would be the expected indicators for the presence of Fe, but only if theeffective temperature is not so high that the population of Fe VII is too much depletedby ionization. Using the plethora of far-ultraviolet transitions lying in the FUSE range,Werner, Rauch, & Kruk (2010) have detected Fe features not only in the cooler PG 1159stars with T eff < FUSE spectra, while in the hotter subset, Fe X is thedetected species and the analysis of abundances are in progress (Werner 2010; Werner et al.1992). Their analyses yielded a solar iron abundance for these stars. These hottest objectsare among the most massive PG 1159 stars (0.71–0.82 M ⊙ ), while those objects revealing thestrongest Fe deficiency are associated with a lower mass range (0.53–0.56 M ⊙ ). Nonetheless,the evolutionary significance, if any, of the presence of Fe in solar abundance in some PG 1159stars versus PG 1159 stars that appear to be Fe-deficient, remains unclear. 10 –It is now widely believed that the hydrogen-deficiency in extremely hot post-AGB starsof spectral class PG 1159 is probably caused by a (very) late helium-shell flash or an AGBfinal thermal pulse (Iben 1984) that consumes the hydrogen envelope, exposing the usually-hidden intershell region. Thus, the photospheric elemental abundances of these white dwarfsoffer insights into the details of nuclear burning and mixing processes in the precursor AGBstars. Werner et al. (2008) compared predicted elemental abundances to those determined byquantitative spectral analyses performed with advanced NLTE model atmospheres. A goodqualitative and quantitative agreement is found for many elemental species (He, C, N, O, Ne,F, Si, Ar), but discrepancies for others (P, S, Fe) point at shortcomings in stellar evolutionmodels for AGB stars. PG 1159 stars appear to be the direct progeny of [WC] Wolf-Rayetstars (Werner, Rauch, & Kruk 2007, 2008), a possibility first suggested by Sion et al. (1985)when the same high excitation O VI absorption features detected in the PG 1159 stars werealso seen in the optical ultraviolet spectra of O VI planetary nebula nuclei.
5. DO White Dwarfs
The DO white dwarfs are hot helium-rich white dwarfs that populate the white dwarfcooling sequence from the hot beginning ( T eff = 120,000 K) down to 45,000 K. The opticalspectra of hot DO stars covering this range of T eff , and newly discovered in the SDSS, aredisplayed in Figure 5.At T eff < M ⊙ ) post-AGB remnant carriesit to effective temperatures as high as 700,000 K on very short time scales (Paczy´nski 1970).However, Werner et al. (1995) have argued that the uhei features in DO white dwarfs cannotbe photospheric because the He II lines would fade completely at such high temperatures.Instead, they proposed an hypothesis of optically thick and hot stellar winds, based on thetriangular shape of the line profile. Furthermore, the lines are blue-shifted, which favors theassumption of an expanding envelope.Related to these uhei DO stars is KPD 0005+5106, one of the hottest know DO stars.Werner et al. (2007, 2008) discovered highly ionized photospheric metals, such as Ne VIII andCa X, requiring extremely high temperature, much higher even than previous analyses thatyielded T eff ∼ T eff = 200,000 K with log g = 6 .
7. The abundances of metals are in therange of 0.7 to 4.3 times solar with an upper limit to any hydrogen present of < . T eff and luminosityit is likely that the chemical abundances are probably affected by a stellar wind. Thus, 12 –diffusion and radiative levitation may not be important factors in controlling the surfaceabundances. Furthermore, Wassermann et al. (2010) found that the chemical abundancesof KPD 0005+5106 most closely resemble the abundances seen in R Coronae Borealis stars.Since the R Cor Bor objects are widely held to be the product of binary mergers (e.g.,Webbink 1984, Han 1998), this may imply that KPB 0005+5106 is itself the product ofsuch a merger and, hence, is the evolved progeny of an R Cor Bor giant. If true, such aconnection would imply that the surface abundances of KPD 0005+5106 are chemical relicsof the progenitor giant and, thus, not controlled by diffusion and radiative levitation. If thisinterpretation is correct, then KPD 0005+5106 would represent a new evolutionary channelproducing DO white dwarfs distinct from the evolutionary channel connecting the PG 1159stars to the DO stars.
6. DB White Dwarfs
The DB stars contain helium to a degree of purity not seen in any other astronomicalobjects. Even at high signal-to-noise ratio, many spectra exhibit only the absorption lines ofHe I. If the DB star has accreted metals from a debris disk, comets, or the interstellar medium,then they are classified DBZ, or DBAZ if hydrogen is present (see below). Some hot DB starsexhibit atomic carbon in the far ultraviolet. The best example is the pulsating hot DB starGD 358 (Sion et al. 1989). It remains unclear if the DBA stars have accreted their hydrogenor if it is primordial and a result of convective mixing. The DB white dwarfs by consensusare the progeny of extremely hydrogen deficient post-AGB stars (e.g., see Althaus et al.2005, 2009, and references therein). The DB cooling sequence extends from the hottest DBstars like GD 358 and PG 0112+122, down to the cooler DB white dwarfs (below 20,000 K),and extending downward to 12,000 K, at which point envelope convection has deepenedsubstantially and the He I lines become undetectable. The distribution by number of thecoolest DO stars to the hottest DB stars (30,000 K to 45,000 K) is interrupted by the DBgap (Liebert et al. 1986). Prior to the SDSS, within this very wide range of temperature, noobjects with H-deficient atmospheres were known to exist.Now, however, the large number of new white dwarfs discovered in the SDSS (Eisenstein et al.2006a) has led to the firm placement of no less than 26 DB stars within the DB gap(H¨ugelmeyer & Dreizler 2009). This raises suspicions that the DB gap was not a real featureof the white dwarf temperature distribution. On the other hand, there is still a deficit ofa factor of 2.5 in the DA/non-DA ratio within the gap (Eisenstein et al. 2006b). However,many other objects whose status is questionable (e.g., DAO, DBA, DAB, masqueradingcomposite DA+DB/DO systems) may alter or eventually erase this deficit. Some of these 13 –objects are found within the gap, while others are seen near the gap edges. There is also thecomplicating factor of circumstellar accretion of hydrogen, and the role played by radiativelevitation and weak winds in this temperature interval.If every DA white dwarf evolving through the DB gap turned into a DB, then thereshould be a significant spike seen in the number of DB stars at the red edge of the DB gap,which is not observed. This is in contrast to the strong signature of convective mixing anddilution that changes (significantly lowers) the DA to non-DA ratio at lower temperatures, T eff < < − M wd ), which amounts to approximately10% of all DA stars cooling through the 45,000 K to 30,000 K interval (Eisenstein et al.2006b). This is contrary to the original contention of Fontaine & Wesemael (1987) that allDA stars should have ultra-thin layers ( < − M wd ). Rather, it appears that the fractionof hot DA white dwarfs that transform into non-DA white dwarfs is on the order of 10% ofall DA stars.The problem of whether a DB gap exists is complicated by the known existence ofseveral peculiar DAB, DBA, or DAO stars believed to lie in the 30,000–45,000 K rangethat (1) show evidence of spectrum variability and/or (2) do not fit atmosphere mod-els, whether homogeneous (completely mixed) in H and He throughout the atmosphereor stratified with the hydrogen all in a very thin upper layer. Several other white dwarfsin or near the DB gap also have peculiar spectra. For example, PG 1210+533, with T eff =45,000 K, exhibits line variability of H, He I, and He II, probably modulated by rotation(Bergeron et al. 1994). Also, GD 323, with T eff = 30,000 K, is a DAB star with a vari-able spectrum that cannot be fit completely successfully by either homogeneous or stratifiedmodel atmospheres (Liebert et al. 1984; Pereira, Bergeron, & Wesemael 2005; Koester et al.2009). Additional examples include HS 0209+0832, with T eff = 36,000 K and a 2% heliumabundance (Jordan et al. 1993), and PG 1603+432, with T eff = 35,000 K and a 1% heliumabundance (Vennes, Dupuis, & Chayer 2004). The existence of these systems adds to themounting evidence that the DB gap is not real, as increasing numbers of He-rich stars arebeing discovered within and near its boundaries. The DBA stars have strong lines of He I and weaker Balmer lines, hydrogen-to-heliumratios (by number) in the range N (H) /N (He) ∼ − to 10 − , and, for the most part, 14 –effective temperature below 20,000 K (Shipman, Green, & Liebert 1987). Hence, they clusterat the low end of the DB temperature distribution. The DBA white dwarfs were previouslythought to comprise roughly 20% of He-rich white dwarfs between 12,500 K and 20,000 K(Shipman et al. 1987). However, more recent surveys have dramatically changed this picture.The SPY project yielded a sample of 71 helium-rich degenerates, of which six were new DBAdiscoveries and 14 were DB stars reclassified as DBA due to the detection of hydrogen lines(Voss et al. 2007). In all, 55% of their SPY sample were DBA stars. This is a factor of almost3 times higher than the fraction of DBA stars first estimated by Shipman et al. (1987).It remains unclear if the DBA stars have accreted their hydrogen or if the small Hmass was originally primordial, diluted by convection, and then floated back to the surfaceas a result of convective mixing. This large fraction of DBA stars, coupled with the totalhydrogen masses estimated for the DBA stars suggests the possibility that DB white dwarfs,as they cool, accrete interstellar hydrogen, thus raising their hydrogen mass to the point atwhich a DBA star appears. This channel for forming DBA white dwarfs was favored if theDB gap was real, because DB white dwarfs could be masquerading as DA stars in the DBgap with only a thin hydrogen layer ( < − M wd ). This layer could be mixed away anddiluted, resulting in a DB star appearing at the cool end of the DB gap. It now appears thatthere is no DB gap. Hence, this constraint on the H-layer mass is no longer relevant. Whileinterstellar accretion of hydrogen cannot be easily dismissed, it may yet prove plausible thatthe hydrogen is accreted from volatile-rich debris or comets. The fact that there are DBAZand DBZA stars with accreted calcium may support this scenario.
7. Hot DQ White Dwarfs
The discovery of hot DQ white dwarfs with carbon-dominated atmospheres (Dufour et al.2007, 2008; also see the chapter in this volume by P. Dufour). in the SDSS Data Release 4sample has raised new questions about white dwarf formation channels. These objects aredistinctly different from the cooler, normal DQ stars, which have helium-dominated atmo-spheres and carbon abundances of log N (C) /N (He) ∼ − by number, with the highest car-bon abundance measured for ordinary DQ stars being log N (C) /N (He) ∼ − . The hot DQsall have effective temperatures between ∼ g = 8, with one object (SDSS J142625.70+575218.4)having a gravity near log g = 9 (Dufour et al. 2008). Their optical spectra contain numer-ous absorption lines of C II, which is the hallmark spectroscopic signature of the hot DQs.Among the strongest transitions of C II are at 4267 ˚A, 4300 ˚A, 4370 ˚A, 4860 ˚A, 6578 ˚A, and 15 –6583 ˚A.Despite extensive searches of the vast SDSS sample, no carbon-dominated DQ stars havebeen found with T eff higher than the hottest hot DQ white dwarf, at 24,000 K. Based uponthis absence of hotter carbon-dominated DQ stars, it is quite possible that these objectsappear helium-dominated at higher temperatures, but with very low mass helium layersthat could be effectively mixed and diluted in the carbon-rich convection zone that formsand deepens due to carbon recombination as the hot DQ star cools (Dufour et al. 2008).However, the helium layer would have to be thin in order to be hidden from spectroscopicdetection as the hot DB transforms into a hot DQ star. Adding to the puzzle posed by thehot DQ stars, Dufour et al. (2009) point out an exceptionally high fraction of hot DQ starswith high magnetic fields ( ∼ ∼
8. Conclusions
It is clear that most of the progress achieved in understanding the competing physi-cal processes in hot white dwarf envelopes, the spectral evolution of hot white dwarfs, andthe identification of the white dwarf progenitor channels has arisen directly from an inter-active combination of synthetic spectral abundances studies via space- and ground-basedspectroscopy, with studies of nuclear astrophysics and thermal instabilities via AGB andpost-AGB stellar evolutionary sequences including mass loss. A major triumph has beenthe successful prediction of surface abundances in hot white dwarfs from born-again thermalpulse models and AGB thermal pulse models. This has led to agreement between the ob-served surface abundances from synthetic spectral analyses of high gravity post-AGB stars(PG 1159 and subluminous Wolf-Rayet planetary nebula nuclei) and the theoretical intershellabundances of their double shell-burning AGB progenitors (Werner & Herwig 2006). Thatpre-white dwarf, post-AGB surface abundances shed light on nuclear astrophysical processesdeep inside the intershell layers of the progenitor AGB star is all the more remarkable.While episodic mass ejection and winds are directly observed in post-AGB stars, includ- 16 –ing planetary nebula nuclei, evolving on the plateau and knee portions of pre-white dwarfevolutionary tracks, it remains unknown whether theoretically-predicted weak wind massloss and ion-selective winds exist along the white dwarf cooling tracks (Unglaub 2008). Yet,along the white dwarf cooling tracks at T eff < T eff > T eff > , Hubeny 1988, Hubey & Lanz 1995; German Astrophysical VirtualObservatory ) to a widening circle of investigators, which should also serve to enhance thequantitative analyses shedding light on the formation and spectral evolution of hot degen-erate stars.It is a pleasure to thank Jay Holberg for useful discussions of hot DA stars and PatrickDufour for discussions of hot DQ stars. I would also like to thank Klaus Werner for providingtemperature data on hot non-DA stars in advance of publication. This work was supported byNSF grant AST1008845, and in part by NSF grant AST807892, both to Villanova University. http://nova.astro.umd.edu
17 –
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This preprint was prepared with the AAS L A TEX macros v5.2.
23 –Table 1. Definition of Primary Spectroscopic Classification SymbolsSpectral Type CharacteristicsDA .............. Only Balmer lines; no He I or metals presentDB .............. He I lines; no H or metals presentDC .............. Continuous spectrum, no lines deeper than 5%in any part of the electromagnetic spectrumDO .............. He II strong; He I or H presentDZ .............. Metal lines only; no H or He linesDQ .............. Carbon features, either atomic or molecularin any part of the electromagnetic spectrumP ........... magnetic white dwarfs with detectable polarizationH ........... magnetic white dwarfs without detectable polarizationX ........... peculiar or unclassifiable spectrumE ........... emission lines are present? ........... uncertain assigned classification; a colon (:) may also be usedV ........... optional symbol to denote variabilityTable 2. Temperature Index Ranges for Hot DA StarsSpectral Type T eff Range (K) 10 × θ eff RangeDA.25 200,000 · · ·
DA.5 100,800 · · ·
DA1 40,320–67,200 1.25–0.75DA1.5 28,800–40,320 1.75–1.25DA2 22,400–28,800 2.25–1.75DA2.5 18,327–22,400 2.75–2.25 24 –Table 3. A Selection of the Hottest DA White Dwarfs ( T eff > a WD Name T eff (K) log g Reference0556 − ∗ +8505 − . +0 . − . Marsh et al. (1997)1248 −
278 70,798 ± . ± .
06 Kidder (1991); Koester et al. (2001)1312 − ∗ ±
906 7 . ± .
04 Kidder (1991); Koester et al. (2001)1312 − ∗ ±
891 7 . ± .
03 Kidder (1991); Koester et al. (2001)0440 − ∗ ±
805 8 .
538 Voss (2006)2244+031 72,000 ± . ± .
08 Homeier et al. (1998)0440 − ∗ ± . ± .
085 Finley, Koester, & Basri (1997)0102 −
185 72,370 ± . ± .
08 Limoges & Bergeron (2010)0556 − ∗ ± . ± .
13 Vennes et al. (1997)1547+015 ∗ ± . ± .
111 Finley, Koester, & Basri (1997)1342+443 ∗ ± . ± .
11 Liebert, Bergeron, & Holberg (2005) † −
049 74,798 ±
944 7 .
475 Voss (2006)1312 − ∗ ±
825 7 . ± .
027 Voss (2006); Koester et al. (2009)1547+015 ∗ ±
561 7 .
612 Voss (2006)0158 −
227 75,758 ± .
386 Voss (2006)0630+200 75,792 ±
751 8 . ± .
050 Finley, Koester, & Basri (1997)1827+778 75,800 ±
610 7 . ± .
03 Homeier et al. (1998)0616 −
084 76,320 ±
200 8 . ± .
15 Vennes (1999)0111 −
381 76,857 ±
746 7 .
367 Voss (2006)1749+717 76,900 ±
550 7 . ± .
03 Homeier et al. (1998)1622+323 77,166 ± . ± .
082 Finley, Koester, & Basri (1997)0441+467 ∗ ± . ± .
14 Bergeron et al. (1994)0939+262 77,300 ± . ± .
06 Bergeron et al. (1994)0229 −
481 77,421 ± . ± .
064 Bragaglia, Renzini, & Bergeron (1995)1547+015 ∗ ± .
49 Liebert, Bergeron, & Holberg (2005) † ∗ ± . ± .
11 Bergeron et al. (1994)1253+378 ∗ ± . ± .
20 Bergeron et al. (1994)1253+378 ∗ ± . ± .
21 Liebert, Bergeron, & Holberg (2005)2146 −
433 81,638 ± .
994 Voss (2006)0441+467 ∗ ± . ± .
13 Bergeron et al. (1994)1305 −
017 85,773 ±
992 7 .
800 Voss (2006)0345+006 86,850 ± . ± .
12 Limoges & Bergeron (2010)1738+669 ∗ +2390 − . +0 . − . Marsh et al. (1997)0500 −
156 94,488 ±
112 7 .
214 Voss (2006)1738+669 ∗ ± . ± .
035 Finley, Koester, & Basri (1997)0615+655 98,000 ± . ± .
15 Homeier et al. (1998)2246+066 98,000 ± . ± .
05 Homeier et al. (1998)
25 –Table 3—Continued
WD Name T eff (K) log g Reference0950+139 108,390 ± . ± . .
38 Liebert, Bergeron, & Holberg (2005)0948+534 136,762 ± .
222 Liebert, Bergeron, & Holberg (2005) † a This list was compiled from the literature with input and advice from J. Holberg(private communication) to select the most reliable temperature estimates obtained fromthe highest quality optical and far-ultraviolet spectra. ∗ Duplicate listing with independent parameter estimates. † Revised parameters provided by J. Holberg (private communication).
Table 4. The Hottest DAO White Dwarfs a Effective Temperature (K)Name Balmer Lines Lyman LinesAbell 7 66,955 99,227HS 0505+0112 63,227 120,000PuWe 1 74,218 109,150RE 0720-318 54,011 54,060Ton 320 63,735 99,007PG 0834+500 56,470 120,000Abell 31 74,726 93,887HS 1136+6646 61,787 120,000Feige 55 53,948 77,514PG 1210+533 46,338 46,226LB2 60,294 87,662HZ 34 75,693 87,004Abell 39 72,451 87,965RE 2013+400 47,610 50,487DeHt5 57,493 59,851GD 561 64,354 75,627 a Data used in this table are from Good et al.(2004). 26 –Table 5. DO White Dwarf Parameters a Metal Abundances b Star T eff (kK) log g C N O Fe NiPG 1034+001 100 7.5 − . − . − . − . < − . − . < − . − . − . − . −
289 70 7.5 − . − . − . < − . − . − . < − . · · · < − . < − . < − . − . < − . · · · · · · HD 149499 B 50 8.0 < − . < − . < − . · · · · · · a Data used in this table are from Dreizler (1999). b Logarithm of number ratios relative to He from homogeneous model atmospherefits. 27 –Fig. 1.— Spectra and model fits of three hot DA white dwarfs (G191-B2B, GD 53, andGD 71) that are used as primary flux calibration standards for
HST . Absorption lines ofthe hydrogen Balmer series are prominent in the spectrum of each star. The fluxes are in f λ units, normalized to have a median value of 1 in the range 3850–5400 ˚A in the top andlower-left panels. The lower-right panel shows the model fit for GD 71 when the continuumshape is rectified. Note that the spectra and model fits are essentially indistinguishablein the plots. From Allende Prieto, Hubeny, & Smith (2009), reproduced by permission ofWiley-Blackwell. 28 –Fig. 2.— Spectra and model fits (smooth lines) of six hot DB white dwarfs, with temper-atures in the range 30,000–45,000 K. The prominent absorption features in each spectrumare lines of He I. From Eisenstein et al. (2006b), reproduced by permission of the AAS. 29 –Fig. 3.— Summary plot of log g versus log T eff showing all analysed PG 1159 stars (filledcircles) and hot DO white dwarfs (unfilled circles), with DO white dwarf cooling tracks byAlthaus et al. (2009). 30 –Fig. 4.— Time sequence of normalized spectra of the central star in Longmore 4, showing(from top to bottom) the appearance and rapid decline of a dramatic emission line phaselinked to a mass ejection event. From Werner et al. (1992), reproduced with permissionc (cid:13) ESO. 31 –Fig. 5.— Normalized optical spectra (grey lines) of DO white dwarfs discovered in the SDSS,along with model atmospheres (black lines), ordered by decreasing effective temperature.From H¨ugelmeyer et al. (2005), reproduced with permission c (cid:13)(cid:13)