ASTRONOMY AND ASTROPHYSICS Spectral analyses of PG 1159 stars: constraints on the GW Virginis pulsations from HST observations

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1 Astron. Astrophys. 334, (1998) ASTRONOMY AND ASTROPHYSICS Spectral analyses of PG 1159 stars: constraints on the GW Virginis pulsations from HST observations S. Dreizler 1,2,3 and U. Heber 3 1 Institut für Astronomie und Astrophysik, Universität Tübingen, Waldhäuser Strasse 64, D Tübingen, Germany 2 Institut für Astronomie und Astrophysik der Universität, D Kiel, Germany 3 Dr.-Remeis-Sternwarte, Universität Erlangen-Nürnberg, Sternwartstrasse 7, D Bamberg, Germany Received 23 December 1997 / Accepted 24 February 1998 Abstract. We present the results of a quantitative analysis of UV and optical spectra of nine PG 1159 stars, very hot hydrogen-deficient [pre-] white dwarfs, by means of line blanketed NLTE model atmospheres. Four programme stars constitute the GW Vir variables, a class of non-radial g-mode pulsators. Precise effective temperatures, carbon, nitrogen and oxygen abundances and spectroscopic masses are used to constrain the GW Vir pulsations. The blue edge of the instability strip is at K (PG ). PG sets the red edge at K, but is also one of the coolest PG 1159 stars known, suggesting that the pulsations are stopped when the transformation of a PG 1159 star into a hot white dwarf occurs by gravitational settling of the metals. Four non-variables are found to lie inside the GW Vir instability strip indicating that an additional parameter determines whether a PG 1159 star pulsates. Abundances of C and O in the pulsating stars appear to be higher than in the non-variables in agreement with the theoretical prediction that the pulsations are driven by cyclic ionization of C and O. The outstanding discovery of our investigation, however, is a strong correlation between the nitrogen abundance and pulsations. All GW Vir stars are nitrogen rich, whereas no nitrogen can be detected in the non-variables except in PG We conjecture that this correlation provides a key for the understanding of the driving mechanism. Comparing their position in the T eff -log g diagram to new evolutionary models we conclude that most programme stars are post-agb stars of rather low mass ( M ) which have lost their entire hydrogen-rich envelope and part of their helium-rich envelope whereas we confirm HS to be an AGB manqué star. The high nitrogen abundance in four stars is a tracer of mixing processes which have led to ingestion and burning of hydrogen during the final helium shell flash. Send offprint requests to: S. Dreizler Based on observations obtained a) with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS ; b) at the German- Spanish Astronomical Center, Calar Alto, operated by the Max-Planck- Institut für Astronomie Heidelberg jointly with the Spanish National Commission for Astronomy; c) at the European Southern Observatory, La Silla, Chile; d) with the International Ultraviolet Explorer (IUE). Key words: stars: abundances stars: atmospheres stars: evolution stars: oscillations stars: AGB and post-agb stars: white dwarfs 1. Introduction The GW Vir instability strip is composed of hot, hydrogen deficient post-agb stars whose prototype PG (=GW Vir) was detected in the Palomar Green Catalog (Green et al. 1986). The peculiar nature of this star became evident when McGraw et al. (1979) took first spectra revealing the high effective temperature from the detection of highly ionized carbon lines. Emission lines at He II 4686 Å and C IV 4660 Å led to speculations about helium mass transfer to a compact star in a new type of an interacting binary. A light curve was taken to verify this hypothesis but instead of the expected flickering a sinusoidal light curve with a period of about eight minutes and an amplitude of 0 ṃ 02 was obtained suggesting that this star defines a new type of low amplitude, non-radial g-mode pulsator. The star PG became also the prototype of a new spectroscopic class, the PG 1159 stars, characterized by a broad absorption trough around 4670 Å composed of He II 4686 Å and several C IV lines. Today, 31 PG 1159 stars are known (see Dreizler et al [DWH], Werner et al. 1996c). Nine of them were found to be variable, five belonging to the group of PNNVs (Planetary Nebulae Nuclei Variables) and four to the GW Vir stars 1. Werner (1992) defined three spectral subtypes of PG 1159 stars (A, E, lge) based on the presence or absence of photospheric emission lines and the width of the absorption wings. A fourth subtype (H) was introduced by Dreizler et al. (1996) to classify stars that show Balmer lines of hydrogen in addition to the characteristic He/C features. These definitions include several central stars of planetary nebula (CSPN) in the PG 1159 class. In fact, most members of the lge (low gravity emission) subclass, such as NGC 246 and K 1 16, are 1 In the literature the pulsators among the luminous PG 1159 stars(lge) are often grouped together with the pulsating [WC] stars as PNNV since their power spectra are very similar. From its spectral and seismological properties HS also belongs to the PNNVs even though no nebula has been detected (Werner et al. 1997b)

2 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars 619 central stars of PN. But CSPN are also found amongst the E (e.g. PG ), A (e.g. Abell 21) and H (e.g. Sh2-68) subclasses. About every other PG 1159 star resides in an old planetary nebula. The quantitative spectral analysis had to await the development of new sophisticated non-lte model atmosphere techniques (Werner 1986). These have been applied to the optical spectra of four PG 1159 stars and effective temperatures, gravities and abundances of C and O have been determined for the first time by Werner et al. (1991 [WHH]). Subsequently, several other PG 1159 stars have been analyzed in a similar way. The results are summarized in the reviews by DWH and Werner et al. (1997a). The quantitative analyses confirmed that PG 1159 stars are amongst the hottest stars known and that their abundance patterns are indeed very peculiar. The PG 1159-phenomenon persists from effective temperatures of K (HS , this work) up to T eff = K (RX J , Werner et al. 1996c), a range considerably wider than thought previously and the region in the HRD overlaps with the DO white dwarfs (Dreizler & Werner 1996). Surface gravities also cover the wide range between log g= Typically, these stars have a surface composition of 33% He, 50% C, and 17% O by mass (PG , WHH). In general, hydrogen is below the detection limit (Werner 1996a). The four PG 1159 stars of type H also display hydrogen lines (therefore also termed hybrid PG 1159 stars; Napiwotzki & Schönberner 1991, Dreizler et al. 1996). Nitrogen is also below the detection limit except for PG with N/He=0.01 (Werner & Heber 1991). The analyses suggested that the PG 1159 stars provide an evolutionary link between the helium- and carbon-rich central stars of planetary nebulae (spectral type [WC]) and the sequence of helium-rich white dwarfs (DO and DB). The link to the [WC] central stars has been strikingly confirmed by observing the dramatic change of spectral type of the central star Lo 4 from a PG 1159 (lge) spectrum to WC2 and back in less than half a year (Werner et al. 1992). The similarity of their surface compositions as determined from spectral analyses of [WC] stars (Koesterke & Hamann 1997, Leuenhagen et al. 1996) corroborated this link. The hot helium-rich white dwarfs (DO) are regarded as successors of the PG 1159 stars. Gravitational settling of the heavier elements, probably retarded by radiative acceleration, turns a PG 1159 star into a DO if no trace of hydrogen is left in the envelope. We emphasize that this trace amount is very difficult to determine since it requires high resolution spectroscopy. With existing spectra we can therefore not exclude that PG 1159 stars turn into hydrogen rich DA white dwarfs (see also Sect. 3.2). In the case of the type H PG 1159 stars, were hydrogen is clearly present, the transition into a DA is more likely. The surface composition of PG 1159 stars cannot be reproduced by canonical evolution calculations, which predict a hydrogen-rich surface during the entire post-agb phase. Therefore, strong mass loss, possibly caused by a late He shell flash (Iben 1984), must be invoked in order to explain the exotic surface abundance pattern which is typical for 3α processed matter of a former, double shell burning, thermally pulsing AGB star which has lost its entire hydrogen and (almost) its helium-rich envelopes. Evidence for such unexpected mass loss events is given by the above mentioned case of Lo 4. Although the PG 1159 stars seem to be exotic, their number is sufficiently large to account for up to 50% of the transition objects from the hottest post-agb phase to the white dwarf stage (Dreizler & Werner 1996). This makes these stars key objects for a complete understanding of post-agb evolution. Even more interesting is the fact that the many observed pulsation modes in the GW Vir stars allow the powerful tools of asteroseismology to be applied as first demonstrated in the case of PG by Winget et al (1991). This provides an independent determination of stellar parameters, which are partly overlapping, partly complementary to the spectroscopical ones, as well as a direct insight into the structure and evolution of these stars. This was the motivation for extensive photometric monitoring of six PG 1159 stars with the Whole Earth Telescope (WET, Nather et al. 1990) as well as in several single site campaigns (PG Winget et al. 1991, PG Fontaine et al. 1991, RX J Vauclair et al. 1993, PG Kawaler et al. 1995, PG Vauclair et al. 1995, O Brien et al. 1996, NGC 246 Ciardullo & Bond 1996). Detailed asteroseismological analyses of WET data (PG , PG ; Kawaler & Bradley 1994, Kawaler et al. 1995) provided many stellar parameters like mass, luminosity, effective temperature, surface composition, thickness of the surface layer, rotational velocities, and upper limits for magnetic field strength, e.g. the mass of 0.59±0.01 M as derived from the g-mode period spacing of PG is in reasonably good agreement with the spectroscopic mass determination (WHH, see also Sect. 4.2). While our existing spectroscopic analyses were sufficient to reduce the huge parameter space of the asteroseismological analysis by providing reliable starting values and to distinguish between multiple solutions (Kawaler et al. 1995), the question of the driving mechanism of the GW Vir pulsation requires more detailed spectroscopic investigations. Even though there is a general agreement among the pulsation theorists (see Gautschy 1997 and references therein) that cyclic ionization of carbon and oxygen in a layer slightly below the photosphere (10 9 M ) causes the instability, details are still unknown. The amount of helium (and hydrogen) in the driving region as well as the stellar radii are critical parameters to match the observed instability strip. Better observational constraints for the limits of the GW Vir instability strip are urgently needed. Up to now the observational limits of the GW Vir instability strip in the HRD are essentially unknown. In fact a puzzle persists: Our spectral analysis of four PG 1159 stars showed them to be two spectroscopic twins (PG = GW Vir and PG ; PG = V817 Her and PG ) with identical atmospheric parameters (to observational limits). However, one star out of each pair pulsates whereas the other does not. Moreover, two of the lowest temperature stars (PG = BB Psc and PG = IR Peg) pulsate while two somewhat hotter objects do not. Hence the physi-

3 620 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars cal parameters which determine whether a star pulsates or not remain obscured. Therefore high quality UV spectra of nine PG 1159 stars (all four known GW Vir pulsators and five non-variable stars 2 ) have been obtained with the Hubble Space Telescope in order to determine the photospheric parameters with much higher precision than from optical data alone, to determine the O abundance for the first time and to reveal the differences between the pulsating and the non-pulsating objects. Precise temperature and O abundance determinations are a prerequisite for any further discussion of the pulsation properties. Especially the oxygen abundance has to be known precisely because of its importance for the driving of the GW Vir pulsations. Unfortunately, only the hottest PG 1159 stars display O lines in their optical spectra from which O abundances have been derived. In contrast to the optical the UV covers several ionization stages of carbon and oxygen so that sensitive temperature indicators become available for analyses. The UV also enables us to determine the oxygen abundance of the cooler ( K, type A) PG 1159 stars by providing suitable spectral lines not available in the optical. All programme stars belong to the high gravity subtype (A or E), PG is a spectroscopic binary (Wesemael et al. 1985) and PG resides in a Planetary Nebula (Jacoby & van de Steene 1995). Seven of them are cool and therefore lack an O abundance determination. The two hot ones ( K) form the spectroscopic twin pair PG /PG discussed above. The former one was observed already in HST cycle 1 (Werner & Heber 1993) and its analysis demonstrated the achievable precision. It will therefore be the guide line for this analysis. In the following sections we describe our observations (Sect. 2) as well as the employed model atmospheres (Sect. 3). Results are discussed in Sect. 4 and summarized in Sect Observations Our HST observations of PG 1159 stars started in cycle 1 with PG using FOS, because the GHRS was not available in that cycle due to the failure of the side 1 power supply. Three more objects (PG , PG , PG ) were scheduled already in cycle 3. However, none of them was observed due to the first service mission. The spectrum of PG was already presented by Werner & Heber (1993), the analysis will however be repeated here with our latest models for the sake of consistency. After demonstrating the high precision obtained in the determination of the effective temperature we re-applied for HST time in cycle 5, this time with an extended target list. After the first refurbishment mission we could then use the GHRS spectrograph which was better suited for our needs. All observations were obtained in the fp-splitfour mode to reduce the fixed-pattern noise. Final spectra were obtained after the standard pipeline extraction. We smoothed 2 photometric monitoring of HS (Vauclair, priv. comm.) and HS (Fontaine, priv. comm.) did not reveal any variability the spectra with a 0.5 Å Gauss profile to reduce the noise which results in a spectral resolution of 0.9 Å for the GHRS and 1.1 Å for FOS data. Details of the observations are listed in Table 1. All optical observations except one have been performed during the past ten years at the German-Spanish Astronomical Center on Calar Alto, Spain, using the 3.5 m telescope equipped with the TWIN or the Boller & Chivens spectrograph. The dispersion varies from 26 Å/mm to 72 Å/mm resulting in a resolution of 1 Åto3.5Å. In the TWIN spectrograph the red and blue channels are separated by a dichroic splitting the light beam at 5500 Å. Observations of the PG 1159 stars found in the Hamburg-Schmidt survey are described in more detail by Heber et al. (1996) and Dreizler et al. (1994a). Spectra of PG , PG , PG , and PG were already presented and analyzed by WHH. MCT was observed by Thomas Rauch (Tübingen, kindly provided for this analysis) at ESO with the ESO Multi Mode Instrument attached to the NTT. A dichroic was used to feed two separate channels. Data reduction was performed in Kiel, Bamberg and Tübingen using standard MIDAS or IRAF routines. For consistency checks we evaluate IUE low- and highresolution data selecting suitable spectra form the Final Archive. Image numbers and exposure times are listed in Table 1. Looking at Fig. 1 it is clear that the resolution and the S/N of the IUE spectra is too low to tackle the problems posed in the introduction. The two important indicators for the effective temperature, the O IV multiplet at 1340 Å and the C III multiplet at 1170 Å are invisible in the IUE spectrum due to the low resolution. A more detailed comparison between HST and IUE low resolution spectra reveals a slight calibration problem. The HST flux redwards of 1380 Å is up to 10% lower than the IUE flux and our model spectra flux. The same result is obtained if we compare the HST spectra with HUT spectra (Kruk & Werner 1998). For the final fits we therefore applied a correction function to our HST data calibrated with PG In Table 2 we list our stellar line identifications from the GHRS spectra. All major features can be attributed to either stellar or interstellar origin (see Verner et al or Morton 1991 for interstellar line lists). As expected, we see several lines of C IIIand C IV as well as O IV,OV, and O VI. Unexpected was the detection of N IV and N V lines in the pulsating PG 1159 stars. The spectral resolution is too low for a separation of stellar and interstellar lines by velocity, all interstellar lines except possibly N V are, however, from low ionization stages (e.g. N I, O I, SiII, SII) and can therefore be clearly distinguished from the high ionization stages of the PG 1159 stars. There are three cases where a safe assignment is more difficult: The contribution of the interstellar N V resonance doublet to the stellar one (see Sect. 4.2.) and the possible identification of a Mg IV line at 1191 Å. According to calculations of Rauch (1997) Mg IV lines could be expected at effective temperatures of PG 1159 stars. However, there should be another Mg IV line at 1307 Å, with comparable line strength and originating from near-by levels, which is not detected. The 1190 Å line is possibly blended by interstellar S III and Si II which would, however, require far too high abundances for these elements if no other contribution

4 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars 621 Table 1. List of HST and IUE observations. N H and E(B-V) are determined from the HST spectra. The flux is given in erg/(cm 2 s 1 Å 1 ) star spectrograph grating t exp/sec IUE image t exp/sec N H/cm 2 E(B-V) 1300 Å HS GHRS G140L SWP HS GHRS G140L SWP MCT GHRS G140L PG GHRS G140L SWP PG FOS G130H SWP SWP SWP SWP SWP SWP SWP SWP PG GHRS G140L SWP PG GHRS G140L SWP SWP PG GHRS G140L SWP PG GHRS G140L SWP Table 2. Identification of stellar spectral lines in HST GHRS spectra; p denotes lines visible only in pulsating, h lines in hot (T eff K), and c lines in cool (T eff < K) PG 1159 stars; b denotes lines possibly blended by interstellar lines originating from low inonization stages (except N V) within 1 Å. λ/å ion rem. λ/å ion rem. λ/å ion rem. λ/å ion rem. λ/å ion rem C IV C IV N IV pc O VI b O VI h O VI h C IV pc O IV O VI h C III c pc O V c pc c He II b O VI h C IV c C IV h O VI h c C IV h C IV O V c h C IV N V pb h pb N IV pc O VI h Mg IV b C III c O VI h h O VI hb C III cb O VI hb C IV h C IV b O VI hb hb h N IV pc C IV C IV pc O V is accounted for. A possible stellar magnesium contribution can therefore only be determined with detailed modeling. This is beyond the scope of this paper. Our strategic O V line at 1371 Å might be blended by interstellar Ni II. The separation of 1.2 Åis large enough to be resolved by the employed GHRS setup. In the final fit of MCT , PG , and PG the weak contribution of the interstellar Ni II line can be seen in the blue wing of the O V line. 3. Spectral analysis 3.1. Model atmospheres and atomic data For our spectral analyses homogeneous, plane-parallel model atmospheres in radiative and hydrostatic equilibrium were applied. Despite of the high gravity the ionization and excitation of the plasma is dominated by the intense radiation field due to the very high effective temperature (WHH). Non-LTE calculations are therefore required to obtain accurate results. We used our non-lte code which is based on the Accelerated Lambda Iteration method (Werner 1986, Dreizler & Werner 1993, Werner & Dreizler 1998) to calculate all atmospheric models presented here. We included the Hummer-Mihalas occupation probability formalism, generalized to non-lte conditions by Hubeny et al. (1994), to ensure the pressure induced dissolution of the atomic levels. The influence on line profiles for hot helium-rich stars is demonstrated by Werner et al. (1995) and Werner (1996b). Without going into details we mention that in general profiles of He II lines become deeper and broader, most prominent is the change in the 4686 Å line. As demonstrated by Dreizler et al. (1994b [DWJH]), the broadening of

5 622 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars relative flux Fig. 1. Comparison of the IUE low-resolution and HST GHRS spectra of HS Top: HST GHRS spectrum; middle: HST GHRS spectrum degraded to the IUE resolution; bottom: IUE low-resolution spectrum. The vertical lines indicate the positions of the important T eff indicators (C III1177 Å; O IV 1340 Å; O V 1371 Å). The strong Lyα geocoronal emission line is omitted for clarity. the C IV resonance line as well as the lines contributing to the trough around 4670 Å has to be accounted for already in the calculation of the atmospheric structure. In order to make full use of the HST spectra we calculated an extended grid of NLTE model atmospheres taking into account most detailed model atoms. The great advantage of the HST spectra is the coverage of several ionization stages of carbon and oxygen. This, however, can only be exploited with very detailed model atoms since all these ionization stages have to be included simultaneously to ensure a consistent analysis of all detectable line transitions in the optical and UV spectra. Altogether we end up with 14 ionization stages, 351 levels and 1140 line transitions for 5 elements taking into account nearly all levels for C III/IV,NV, and O IV/V/VI listed by Bashkin & Stoner (1975). Atomic models are from WHH and DWJH partly updated with Opacity Project data (Seaton 1987). Using these model atoms we calculated an extensive grid of 90 model atmospheres varying the effective temperature, surface gravity and surface composition. Due to the detailed model atoms included in the calculations an enormous amount of computer time was required. All models were calculated on CRAY Y-MP machines of the Computer Center of the University Kiel each one consuming roughly 5 days of CPU time. This effort is due to slow convergence especially caused by the detailed oxygen model atom. Finally, the emergent spectrum of each model was computed on Alpha workstations of the University Tübingen taking into account the best available broadening theories. Stark broadening of the C IV lines is the main source of uncertainty in our analysis since the observed lines are within the transition range of linear and quadratic Stark broadening regimes which is difficult to handle theoretically. We therefore still have to rely on the approximation introduced by WHH for these lines. Keeping this problem in mind it is not surprising that the profiles of the C IV lines around 1350 Å and 1315 Å can not be reproduced exactly (see Figs. 2a d and 4). For the broadening of He II lines we used the tables of Schöning & Butler (1989) which are based on the Unified Theory of Vidal et al. (1970) Spectral analysis of PG 1159 stars Determination of the stellar parameters is performed by comparison of the observed spectrum with theoretical spectra. The fit procedure is demonstrated in Figs. 2a d and 3a and b for the pulsator PG because of its importance for asteroseismology. Its spectral analysis is hampered, though, by a cool companion (Wesemael et al. 1985) whose continuum contribution has been subtracted here. The emission cores in the Hα and Hβ lines (see Figs. 3a and b and 5) arising from the cool companion, however, cannot be corrected for (see also Paunzen et al. 1998). The final fits are displayed in Figs. 4 to 6, the results of the analyses are summarized in Table 3. The fitting procedure was performed by eye starting with the parameters determined from previous optical analyses (DWH, WHH). First we determined the interstellar column density of hydrogen by fitting the line wings of Lyα and using the standard Seaton law (1979) the reddening by fitting the slope of the UV continuum which is, within the error limits for the effective temperature, only dependent on E(B-V). Varying the model parameters we then derived the best fitting model. The determination of the effective temperature is based on the strengths of the C III 1170 Å, O IV 1340 Å, and O V 1371 Å lines. All three temperature sensitive lines agree very well in the final model providing the expected low error range of ±5%. An additional O V line at 1418 Å confirms the results. Determination of the surface gravity is much more difficult since no particularly gravity sensitive line exists in the range of our HST spectra. Gravity determination is further hampered by the problem of line broadening (see above). We therefore calculated models only with two different surface gravities (log g=7.0 and 7.5) representing the typical (high gravity) PG 1159 star surface gravities. From HST spectra alone only a preference for one of the two gravities can be derived. No interpolation between the two values was therefore applied. Finally the chemical composition is determined by fitting the line strengths of the element in question (C, N and O). During the analysis it turned out that we could only determine the C/O ratio rather than their absolute values if the carbon abundance exceeds C/He=0.3 by number. Since the UV continuum is dominated by C and O opacities the line strengths of C and O lines are insensitive to changes in C and O abundances as long as the C/O ratio is kept fixed. Optical spectra are needed in addition to derive absolute values (see below). Finally we derived N abundances from the resonance doublet at N V 1240 Å for PG , PG , PG , and PG or upper limits for the others which show only very weak N V 1240 Å lines. Since these weak lines could also include a contribution of interstellar matter these values are only upper limits. In the case of PG , PG , and PG the derived abundance is consistent with weak N IV lines at 1275 Å and 1296 Å.

6 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars 623 a T eff 2 b T eff 2 80 kk 90 kk relative flux 1 80 kk kk 90 kk 110 kk CIII 100 kk OIV OV kk c T eff =90kK log g d T eff =100kK, log g=7.5 C/O relative flux 2 2 C/He CIV CIV CIV OIV CIV OV Fig. 2a d. Sample fit of PG Observations are plotted with thin, theoretical spectra with thick lines. Stellar lines are marked at the bottom, interstellar lines at the top. The temperature sensitive C III, OIV and O V lines allow a determination with 5% accuracy (a, b). The gravity determination is hampered by the poorly known broadening of the C IV lines. A systematic error lower than 0.3 dex is therefore impossible to achieve (c). The C/O ratio can be determined with an uncertainty of a factor of three (d top two). The C/He abundance can only be determined in the less carbon rich stars with an uncertainty of a factor of two (d lower three). In the more carbon rich stars the derived C/He abundances from the HST spectra are only lower limits. Optical spectra are needed to improve the results (Fig. 3a and b). The parameters of the final models are then checked by a comparison with optical and IUE spectra. The IUE spectra cover the C IV 1550 Å resonance doublet and the He II 1640 Å line, which are both outside the GHRS range. In all cases except for PG both lines are well reproduced with the HST parameters. A further confirmation of our results comes from a fit to HUT spectra of PG and PG with our HST parameters (Kruk & Werner 1998). With the effective temperature fixed by the UV analysis we hoped to improve the gravity and the He abundance determined from optical spectra. Too high a carbon abundance results in too shallow optical He II lines. Thereby the C/He ratio was

7 624 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars a log g b C/He relative flux Fig. 3a and b. Sample fit of He II 4860 Å in the optical spectrum of PG Observations are plotted with thin, theoretical spectra with thick lines. The emission reversal in the He II line core is due to Hβ emission from the cool companion (see Paunzen et al. 1998). We subtracted its continuum contribution but could not correct for the emission components. Gravity determination is also difficult in the optical. Spectra with higher signal/noise and resolution are required to reduce the error bar (a). The He II lines are most sensitive to changes in the C/He ratio. While a higher ratio than C/He=0.3 could not be excluded from the HST spectra (Fig. 2a d), the He II lines become too weak in this case (b). constrained to within a factor of two. The gravity determination, however, could not be improved significantly. Although the broadening of the He II lines is more reliable the precision of the gravity determination is limited by the available spectra and is not better than from the UV metal lines ( ±0.5 dex, see Fig. 3a and b). At least, we find no discrepancy between both gravity indicators, lending some support to our approximate treatment of CIV/OVI line broadening. The N V 4604/4620 Å doublet provides an independent check of the N abundance. In the case of PG the upper limit could be reduced significantly compared to the HST value. In the case of PG the determination of the N abundance is inconclusive since the doublet is just in the transition from absorption to emission in the parameter range of PG (Fig. 7). While the resonance line in the HST spectrum indicates an N abundance of 1% the optical N V lines are much too strong at the parameters derived previously (T eff = K, log g=7.0). However, slight changes within the error limits weaken the emission line of the /7.0 model to an undetectable strength in the /7.5 model, in accordance with the observation. A slightly lower effective temperature would then turn the line into an absorption line. Due to this sensitive dependence of the N abundance on other parameters the optical spectrum of PG does not constrain it further. Because of the high effective temperature of PG the resonance doublet in the UV is weak and an interstellar origin can not be excluded. However, we have two independent hints that it is more likely of stellar origin. First, the N V lines in PG are significantly stronger than in PG which has otherwise very similar stellar parameters but a higher interstellar hydrogen column density. Second, Table 3. Results of the analyses. Top: non-pulsating, bottom: pulsating PG 1159 stars. Abundances are given in number ratios. star T eff /kk log g C/He O/He N/He PG <10 4 PG <10 5 HS <10 5 MCT <10 5 HS <10 5 PG PG PG PG the IUE high-resolution spectra of PG allow a separation of the stellar and interstellar lines (Liebert et al. 1989). While the interstellar lines have a mean velocity of km/s the N V lines have 33.8 km/s in reasonable agreement with other stellar features (O V 1371 Å, C IV 1548/1550 Å). The separation of 60 km/s, however, is only marginally above the limit of the spectral resolution of the IUE Echelle spectrum. For a conclusive determination we therefore need additional spectra in another wavelength range to corroborate our result. One possibility is the FUV where further N V lines can be found. An existing HUT spectrum of PG (Kruk & Werner 1998) indicates the presence of N at the 1% level, however, the resolution is not sufficient for a safe statement. The problem can only be solved after the launch of FUSE (Far Ultraviolet Spectroscopic Explorer) which will obtain a high resolution FUV spectrum of PG

8 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars 625 NV PG T eff log g C/He C/O N/He kk < 10-2 PG kk < 10-5 HS kk < 10-5 MCT kk < 10-5 relative flux 2 HS PG kk < kk PG PG kk PG kk kk CIII NV OIV OV Fig. 4. Final fits of de-reddened HST spectra from PG 1159 stars. Observation thin, theoretical spectra thick lines. Top five: non-pulsating, bottom four: pulsating PG 1159 stars. Note the strong N V doublet at 1238/1242 Å in the cool pulsators indicated by vertical lines at the bottom. The nitrogen abundances in the stable PG 1159 stars are upper limits as derived from the HST spectra. In the case of PG a tighter limit can be derived from the optical spectrum. Also indicated at the bottom are the strategic lines for T eff determinations, the interstellar lines originating from low ionization stages (e.g. N I, OI, SiII, SII) are marked at the top. A contribution of high ionization interstellar lines is possible only in the case of the N V resonance doublet (see text). Abundances are given as number ratios. For identification of the spectral lines and possible interstellar contamination see Table 2. A progress in the determination of the hydrogen abundance is only possible with high resolution spectra ( 0.1 Å). Such spectra are not yet available for all programme stars except PG Werner (1996a) could derive an upper limit of 5% in this case. It should be noted that a hydrogen abundance of 5% would still be sufficient to transform a PG 1159 star into DA white dwarf during the further evolution. As is obvious from Figs. 4 to 6 the fits of the observed UV and optical spectra are excellent in most cases. Slight discrepancies occur only in the C IV 1350 Å and 1315 Å lines as well as in the He II/Hα line in few cases. The former problem is probably due to inaccurately known broadening theory (see Sect. 3.1). Several lines in the UV spectra are not matched by our theoretical models. All these lines are of interstellar origin. The resolution of the employed gratings, however, does not allow a detailed analysis of the interstellar clouds, except for the hydrogen column density (see Table 1 for results).

9 626 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars T eff log g C/He C/O N/He PG kk < PG HS kk < kk <10-5 MCT kk <10-5 relative flux HS PG kk < kk PG kk PG kk PG kk Fig. 5. Final fits to normalized optical spectra. Observation thin, theoretical spectra thick lines. Top five: non-pulsating, bottom four: pulsating PG 1159 stars Results The resulting atmospheric parameters and C, N and O abundances are summarized in Table 3. Effective temperatures range from K to K at gravities of either 7.0 or 7.5. The atmospheric parameters of PG derived by WHH are confirmed. The former twins (PG /PG and PG /PG ), however, are disrupted by the new analysis, indicating that PG is slightly hotter and has a higher gravity than PG corroborating earlier results of Werner et al. (1996b) which were based on the analysis of an EUVE spectrum. In the other twin pair PG turns out to be hotter than found by WHH, whereas PG is cooler and has a higher gravity. HS is also hotter than found by DWJH. The changes in T eff and log g are within the error limits of previous analyses (10-15%, 0.5 dex, respectively). While the programme stars (except HS , see Sect. 4.4) have very similar C and O abundances, surprisingly four stars are N rich (all with the same abundance of 1% by number relative to He) while no nitrogen can be detected in the others indicating that N/He is below the solar abundance. The derived C and O abundances agree also with previous results to within their error limits. Changes are due to (i) the modified effective temperatures and gravities and (ii) the improvements in the model atmospheres described in Sect. 3.1.

10 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars PG PG HS relative flux 3 2 HS PG PG PG PG Fig. 6. Final model spectra from HST fitting compared to de-reddened IUE low resolution spectra. In addition to the HST spectra IUE covers the C IV lines at 1550 Å and 1585 Å and the He II line 1640 Å. Top four: non-pulsating, bottom four pulsating PG 1159 stars. For line identification see WHH. 4. Discussion We want to discuss our results in the context of asteroseismology and post-agb evolution. Our analyses provide several interesting implications for both fields Effective temperatures and the boundary of the GW Vir instability strip The main motivation for the HST observations was a precise determination of T eff. The better access of T eff from UV spectra now provides a reliably defined instability strip in the HRD (Fig. 8). The pair PG and PG defines the blue edge of the GW Vir instability strip since PG is variable and PG is stable. Due to a better determination of T eff these stars are separated by K now. However, from preliminary analyses NGC 650 (the stable star next to PG in Fig. 8) has parameters similar to Fig. 7. Determination of the nitrogen abundance of PG Comparison of the optical, HST and HUT spectra with theoretical spectra. Positions of the nitrogen lines are marked in the HST and HUT spectrum. While the HST and HUT data can be fitted with the K/log=7.0 model, the N V lines in the optical are too strong. A reduction of the effective temperature to K and an increase of log g to 7.5 provides a good fit, however. See text for detailed discussion. PG but is stable. A more elaborated analysis of NGC 650, however, is required for more detailed conclusions. So, the puzzle remains that stable and instable PG 1159 stars can have very similar effective temperatures and gravities. An upcoming detailed analysis of NGC 650 might reveal differences between these two stars. The red edge seems to be in

11 628 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars Table 4. Comparison between spectroscopic and asteroseismologic masses in M (top four). Spectroscopic masses are derived from the comparison with tracks of O Brien & Kawaler (in prep.). In the case of PG and HS the extrapolation is quite large and the masses can only be guessed. HS is probably not an AGB descendant as discussed in DWJH. The error in the spectroscopic mass determination is 0.1 M. star M spec M ast comment PG Kawaler & Bradley 1994, WET PG Kawaler et al. 1995, WET PG Fontaine et al PG O Brien et al PG PG HS MCT accordance with the low temperature limit of the spectroscopic class of PG 1159 stars. HS and PG are now defining the cool end of PG 1159 stars, the former is stable the latter is pulsating. This clearly indicates that the pulsations stop due to abundance evolution effects: Around the position of HS and PG the transition from PG 1159 stars towards white dwarfs seems to be completed. Heavier elements are then removed from the outer envelope due to gravitational settling making the star appear as a hot DO (or DA if a sufficient amount of H remained in the envelope) white dwarf and destroying the C/O driving mechanism Masses and the width of the GW Vir strip One important point is that the instability strip becomes considerably narrower compared to previous results (Werner et al. 1996b). In the regime of the GW Vir stars the instability seems to be confined to a very small mass range around 0.55 M in contrast to the PNNV instability strip which is much wider. Predictions of the strip position from pulsational models are still not satisfactory as mentioned in Sect They either require a steep abundance gradient to obtain a pure C and O mixture in the driving region (Bradley & Dziembowski 1996) or predict much less luminous or cooler GW Vir stars (Gautschy 1997, Fig. 8). The latter problem is probably due to the initial conditions of the evolutionary sequences of Gautschy. It is therefore important to compare our results with asteroseismological analyses (which are independent from the models concerning the position of the GW Vir strip). One common parameter is the stellar mass which can be determined with high precision from asteroseismology. In order to circumvent systematic effects we derive our spectroscopic masses from a comparison with the evolutionary tracks also applied for the asteroseismologic analyses (O Brien & Kawaler in prep., calculated with input physics similar to Dehner & Kawaler 1995; tracks are kindly provided by these authors). This is necessary as can be clearly seen from Fig. 8 where we included different evolutionary tracks. Previously we used the ones of Wood & Faulkner (1986), which have however a hydrogen rich surface. There is a systematic offset of 0.03 M compared to the tracks of O Brien & Kawaler which are hydrogen deficient and have a surface composition similar to our spectroscopic results. The main uncertainty of the spectroscopic mass determination is the uncertainty of the surface gravity. An error of 0.3 dex translates into an error of 0.1 M. In Table 4 we compare our spectroscopic masses with those from asteroseismology. In the two cases of WET observations the values are in reasonable agreement (PG , PG ) within our error limits, in the two cases of single site observations (PG , PG ) the spectroscopic mass is significantly lower. Upcoming analyses of the latest WET runs will hopefully clarify this discrepancy. Another parameter which can be compared is the effective temperature. This in very good agreement between spectroscopy and asteroseismology, e.g. for PG K (this paper) versus K (Kawaler & Bradley 1994). Asteroseismology also corroborates the surface composition determined from spectroscopy (Kawaler & Bradley 1994). Together, this can be regarded as good mutual confirmation of both approaches. Future spectroscopic work will concentrate on more precise gravity determinations. Finally, it is worthwhile to note that PG is more massive (0.65 M ) than any other programme star ( M ). It is also the only star in the sample known to reside in a PN Abundances and driving mechanism Nitrogen abundance A surprising result of our analysis is the detection of nitrogen in four of the programme stars. Only one PG 1159 star with detectable amount of N was known previously (PG , Werner & Heber 1991). In the other stars the optical N lines were below the detection limit of older spectra. The result is also extremely interesting in view of the driving mechanism of the GW Vir pulsations. In all analyzed stars of this sample the existence of N is strongly correlated with pulsations. All stable PG 1159 stars of the sample have a N/He abundance below while pulsators have a N/He abundance at the 1% level. PG seems to fit into this correlation also. It needs, however, confirmation by future observations. Since our sample contains all GW Vir variables we conclude that a high N abundance is necessary to drive oscillations in this parameter range or directly traces the necessary driving process. Compared to the high C and O abundances, the N abundance (1%) in the pulsating stars is relatively low. Whether this is too low to drive pulsations remains to be seen. It should be noted, though, that the derivatives of the opacity with respect to the local temperature and densities are important for the driving, not the total opacity. The other PG 1159 star with nitrogen, PG , is somewhat hotter (T eff = K) and of lower gravity (log g=6.5) than PG Despite of its nitrogen-richness no variabil-

12 S. Dreizler & U. Heber: Spectral analyses of PG 1159 stars 629 Lo 4 NGC 246 log g (cgs) RX J RX J PG PG NGC 650 Lo 3 PG H PG K 1-16 HS PG log T eff /K Jn 1 Abell 21 VV 47 HS IW 1 HS HS MCT PG PG PG Fig. 8. T eff -log g diagram with observed locations of PG 1159 stars. Different symbols are used to distinguish pulsators and non-pulsators (squares: pulsators, circles: non-pulsators, triangles: no photometric data available). Big symbols represent our programme stars. The parameters for the other PG 1159 stars are taken from Werner et al. (1997a). Post-AGB stellar evolutionary tracks of Wood & Faulkner (1986, He burners, mass loss type B) are plotted with long dashed lines: Masses are from right to left 0.6, 0.7 and 0.76 M. Solid lines are used for tracks of O Brien & Kawaler (in prep.). Models of 0.550, 0.573, 0.600, and M have a helium and carbon rich surface similar to the observed atmospheric abundances. Evolutionary tracks which have started from homogeneous He main sequence with 0.57, 0.63, and 0.70 M (Gautschy 1997) are marked with short dashed lines, thick parts denote phases where pulsational instability occurs. ity has been detected (Grauer et al. 1987). Admittedly this is at odds with our proposed correlation. It is worthwhile to notice that previously Vauclair (1990) suspected nitrogen to be responsible for the GW Vir pulsations. He found from diffusion calculations that N should dominate in the driving region. Since however, the abundance pattern of PG 1159 stars is impossible to explain in the context of diffusion and radiative levitation and due to the fact that only one PG 1159 star with atmospheric N abundance was known at that time this explanation was abandoned again. Does a correlation of the N abundance with pulsations exist in the PNNVs? No N could be detected in any of the PG 1159 stars of lge subclass (e.g. RX J , Werner et al. 1996a) so far, whether pulsating or not. In the case of the [WC] stars some show nitrogen and some don t with no correlation to pulsations. Additionally, the two stable PG 1159-[WC] transition objects (A 30 and A 78) show nitrogen. The PNNV pulsators, though, differ from the GW Vir stars with respect to their stellar winds. Fast winds ( 3000 km/s) have been detected in UV spectra of the central stars of NGC 246 (Heap, 1982) and K 1-16 (Patriarchi & Perinotto, 1996) and are likely to be present in the other PG 1159 lge stars as well. The existence of the stellar wind might strongly influence the pulsations. Both the PG 1159 stars among the PNNVs and the pulsating [WC] stars have pulsation frequencies around 1 mhz while GW Vir stars have frequencies between 2 and 2.5 mhz. Furthermore the frequency pattern of the former group is variable (Steven Kawaler coined the term variable variables ) while it is much more stable in the case of the latter group. In the light of the results regarding the N abundances we conjecture that also the driving mechanism is (slightly) different between the PNNVs and GW Vir stars despite of otherwise very similar spectroscopic and evolutionary properties Carbon and oxygen abundances Another important point is that we could derive O abundances for the cooler (T eff K) PG 1159 stars, which was impossible from previous optical analyses. It was one of the main motivations for the HST proposal to find differences in the C and O abundance between pulsators and non-pulsators since these elements are believed to cause the instabilities due to cyclic ionization (e.g. Starrfield et al. 1984). Indeed, the non-variable HS has considerably less C and O than the GW Vir stars. The other non-variables have also C, O and O/C on average lower than the pulsating stars, but the difference is small (less than a factor of two) and therefore only marginally significant. Hence, we conclude that a correlation between pulsation and C and O abundances might exist, however, considerably weaker than in the case of the N abundance.