Free Access
Issue
A&A
Volume 598, February 2017
Article Number A135
Number of page(s) 30
Section Stellar atmospheres
DOI https://doi.org/10.1051/0004-6361/201629869
Published online 14 February 2017

© ESO, 2017

1. Introduction

The white dwarf (WD) RX J0503.9–2854 (henceforth RE 0503–289, WD 0501−289McCook & Sion 1999a,b) was discovered in the ROSAT (ROentgen SATellite) wide field camera all-sky survey of extreme-ultraviolet (EUV) sources (Pounds et al. 1993). Barstow et al. (1993) reported its discovery by the Extreme Ultraviolet Explorer (EUVE), and identified it with a peculiar He-rich DO-type WD, namely MCT 0501−2858 in the Montreal-Cambridge-Tololo survey of southern hemisphere blue stars (Demers et al. 1986). They found that RE 0503–289 is located in a direction with very low density of the interstellar medium (ISM). In the line-of-sight (LOS) toward RE 0503–289, Vennes et al. (1994) measured a column density of log (NH i/ cm-2) = 17.75 − 18.00 using EUVE photometry data. Rauch et al. (2016a) resolved at least two ISM components in the LOS toward RE 0503–289 based on high-resolution and high signal-to-noise ultraviolet (UV) spectroscopy performed by Far Ultraviolet Spectroscopic Explorer (FUSE) and Hubble Space Telescope/Space Telescope Imaging Spectrograph (HST/STIS) and measured a very low (EBV = 0.015 ± 0.002) interstellar reddening.

Table 1

History of Teff and log g determinations (cf., Müller-Ringat 2013).

The almost negligible contamination by ISM line absorption allows us to identify even weak lines of many species from so far He up to trans-iron elements as heavy as Ba (Table 2). For reliable abundance analyses of these elements, a precise Teff and log g determination is a prerequisite to keep error propagation as small as possible. An initial constraint of Teff = 60 000 − 70 000 Kwas given by Vennes et al. (1994) from EUV photometry. The first spectral analysis by means of non-local thermodynamic equilibrium (NLTE) stellar atmosphere models considering opacities of H, He, and C was published by Barstow et al. (1994). They found Teff = 60 000 − 80 000 Kand log (g/ cm / s2) = 7.5 − 8.0. Dreizler & Werner (1996) used ultraviolet (UV) spectra in addition and NLTE model atmospheres and determined Teff = 70 000 ± 5000 Kand log g = 7.5 ± 0.5. Recently, Rauch et al. (2016a) analyzed optical and ultraviolet (FUSE and HST/STIS) spectra and significantly reduced the error limits to ± 2000 K and ± 0.1, respectively. Table 1 summarizes previous analyses.

2. Observations

In this paper, we used the observed spectra that are briefly described in the following. If they are compared to synthetic spectra, the latter are convolved with Gaussians to model the respective instrument’s resolution.

Extreme ultraviolet:

observations by the EUVE observatory were performed using the short-wavelength (70 Å <λ< 190 Å), the medium-wavelength (140 Å <λ< 380 Å), and the long-wavelength (280 Å <λ< 760 Å) spectrometers with a resolving power of R ≈ 300. Details of the data reduction are given by Dupuis et al. (1995).

Far ultraviolet:

spectra (910 Å <λ< 1190 Å, R ≈ 20 000) were obtained with FUSE. Their data IDs are M1123601 (2000-12-04), M1124201 (2001-02-02), and P2041601 (2000-12-05). The spectra were shifted to rest wavelengths and co-added. For details see Werner et al. (2012b).

Ultraviolet:

spectroscopy was performed with HST/STIS on 2014-08-14. Two observations with grating E140M (1144 Å <λ< 1709 Å, R ≈ 45 800) and two observations with grating E230M (1690 Å <λ< 2366 Å, 2277 Å <λ< 3073 Å, R ≈ 30 000) were co-added. These observations are retrievable from the Barbara A. Mikulski Archive for Space Telescopes (MAST).

Optical:

spectra (3290 Å <λ< 4524 Å, 4604 Å <λ< 5609 Å, 5673 Å <λ< 6641 Å) were obtained on 2000-09-09 and 2001-04-08 in the framework of the Supernova Ia Progenitor Survey project (SPY, Napiwotzki et al. 2001, 2003). The Ultraviolet and Visual Echelle Spectrograph (UVES) attached to the Very Large Telescope (VLT) located at the European Southern Observatory (ESO) on Cerro Paranal in Chile was employed to achieve a resolution of about 0.2 Å. In addition, we use a spectrum taken with the Echelle Multi Mode Instrument (EMMI) attached to the New Technology Telescope (NTT; 1992-01, 4094 Å <λ< 4994 Å, resolution of about 3.0 Å).

Near infrared:

spectroscopy (9500 Å <λ< 13 420 Å, R ≈ 950) was performed on 2003-12-10 using the Son-of-Isaac (SofI) instrument at the NTT. The spectrum used here was digitized with Dexter1 from Fig. 1 in Dobbie et al. (2005).

Table 2

Photospheric abundances (mass fraction) of RE 0503–289.

3. Model atmospheres and atomic data

The stellar model atmospheres used for this paper were calculated with our Tübingen NLTE Model Atmosphere Package (TMAP2, Werner et al. 2003, 2012a). They assume plane-parallel geometry, are chemically homogeneous, and in hydrostatic and radiative equilibrium. An adaptation is the New Generation Radiative Transport (NGRT) code (Dreizler & Wolff 1999; Schuh et al. 2002) that can consider diffusion in addition to calculate stratified stellar atmospheres.

The Tübingen Model Atom Database (TMAD3) provides ready-to-use model atoms in TMAP format for many species up to Ba. TMAD has been constructed as part of the Tübingen contribution to the German Astrophysical Virtual Observatory (GAVO4).

Werner et al. (2012b) discovered lines of trans-iron elements, namely Ga (atomic number Z = 31), Ge (32), As (33), Se (34), Kr (36), Mo (42), Sn (50), Te (52), I (53), and Xe (54), in the FUSE spectrum of RE 0503–289. For precise abundance determinations of these species, reliable atomic data is mandatory. For example, reliable transition probabilities are required, not only for lines that are identified in the observation but for the complete model atoms that are considered in the model-atmosphere calculations. Due to the lack of such data, Werner et al. (2012b) were restricted to abundance determinations of Kr and Xe only.

We initiated the calculation of new transition probabilities that were then used to determine the abundance of the respective element. Table 3 gives an overview of the so far calculated data. To provide easy access to this data, the registered Tübingen Oscillator Strengths Service (TOSS) has been created within the GAVO project.

To construct model atoms for the use within TMAP, the elements given in Table 3 require the calculation of so-called super levels and super lines with our Iron Opacity and Interface (IrOnIc, Rauch & Deetjen 2003) due to the very high number of atomic levels and lines. We transferred the TOSS data into Kurucz’s data format5 that can be read by IrOnIc.

4. Radial velocity and gravitational redshift

To shift the observation to rest wavelength, we determined the radial velocity vrad of RE 0503–289 from FUSE and HST/STIS spectra. To measure the wavelengths of the line centers, we used IRAF6 to fit Gaussians to the line profiles. In total, we evaluated 100 lines in the FUSE wavelength range and 103 lines in the STIS wavelength range (Fig. 1). The averages are and . We adopted the mean value of . From this value, the gravitational redshift z has to be subtracted. To calculate z and the respective radial velocity, we created the GAVO tool Tübingen Gravitational REDshift calculator (TGRED, Fig. C.2). For RE 0503–289, we derive . The true radial velocity is then .

Table 3

Newly calculated transition probabilities.

thumbnail Fig. 1

Determination of vrad from individual lines in the FUSE (top panel) and HST/STIS observations (bottom). The full horizontal lines indicate the average vrad for FUSE and HST/STIS, respectively. The dashed lines show the 1σ error.

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5. Line identification

To unambiguously identify lines in our observed spectra (Sect. 2), we used the best synthetic model of Rauch et al. (2016a) and calculated additional spectra with oscillator strengths set to zero for individual elements. This allows to find weak lines, even if they are blended by stronger lines. The detection limit is an equivalent width of Wλ = 2 mÅ. Table 4 shows the total numbers of lines identified in the four wavelength ranges and the numbers of lines that were suited to determine Wλ and vrad. The current line lists are presented in Tables A.1–A.5, a regularly updated version is available online7.

Table 4

Statistics of the identified (in brackets: newly identified in this paper) and unidentified lines in the observed spectra.

6. Visualization and online line list

In the framework of the Tübingen GAVO project, we have developed the registered Tübingen VISisualization tool (TVIS) that allows the user to plot any data in an easy way on the WWW. The plotter itself is written in HTML5 and Javascript. To strongly increase the security of this web application, no Flash or Java is necessary to use it, meaning that TVIS will even work when Flash is dead and Java applets are blocked by the browsers.

The comparison of our best model spectra with the available observation of RE 0503–289 in the EUV, FUV, NUV, and optical wavelength ranges was realized with TVIS and is shown online8. Figures B.1 to B.3 show the FUV to optical range.

thumbnail Fig. 2

Determination of EBV. Top: reddening with EBV = 0.00026 applied to our synthetic spectrum in the wavelength range from the EUV to the NIR. Bottom: same like top panel, for EBV = 0.00016 (blue), 0.00026 (red), and 0.00036 (green) in the EUV wavelength range. Prominent lines are marked.

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thumbnail Fig. 3

Comparison of the EUVE observation (gray line in both panels) with our models. Top panel: two models with Teff = 70 000 K (red) and Teff = 66 000 K (blue). Identified photospheric lines are marked at the top. Bottom panel: three models with Teff = 70 000 K. Red, thick line: model from the top panel, red, dashed line: model with 200 times increased Ni abundance, blue, thin line: model that considered only opacities of He, C, N, O, Fe, and Ni.

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7. Is there still an EUV problem in RE 0503–289?

To analyze the EUVE observation, Barstow et al. (1995) used NLTE model atmospheres that were calculated with the code that is nowadays called TMAP. A synthetic spectrum (scaled to match the observed EUV flux) that was calculated from a model with Teff = 70 000 K, log g = 7.0, C / He = 1%, and N / He = 0.01% (the latter being number ratios) reproduced well the observation. A major problem arose, however, from the fact that the model flux (reddened and interstellar neutral hydrogen absorption considered) in the wavelength range 228 Å <λ< 400 Å was about an order of magnitude higher than observed. Only models with Teff<65 000 K produced an acceptable fit. He iλ 5875.62 Å (2p 3Po–3d 3D) in the optical wavelength range (e.g., in spectra taken with the TWIN spectrograph at the Calar Alto observatory and in SPY spectra, Dreizler & Werner 1996; Rauch et al. 2016a) establishes a stringent constraint of Teff = 70 000 ± 2000 K.

Werner et al. (2001) calculated TMAP models that were composed of He, C, O, and the iron-group elements (Ca–Ni). Interstellar He i absorption was applied in addition to that of H i. The flux discrepancy was reduced (model flux three times higher than observed) but the basic problem, finding an agreement at Teff = 70 000 K, was not solved.

Müller-Ringat (2013) created the Tübingen EUV absorption tool (TEUV, Fig. C.1), that corrects synthetic stellar fluxes for ISM absorption for λ< 911 Å. Presently, only radiative bound-free absorption of the lowest ionization states of H, He, C, N, and O is simulated. Opacity Project data (Seaton et al. 1994) is used for the photoionization cross-sections. These consider, for example, autoionization features. For this paper, Si has been added to TEUV. Two interstellar components with different radial and turbulent velocities, temperatures, and column densities can be considered. Müller-Ringat (2013) calculated TMAP models (Teff = 70 000 K, log g = 7.5) that included He, C, N, O, and the iron-group elements. Although Kurucz’s line lists were strongly extended in 2009 (Kurucz 2009, 2011), and about a factor of ten more iron-group lines were considered, the EUV model flux was about twice as high as that observed. To match the observed EUV flux, Teff had to be reduced to ≲ 65 000 K.

Rauch et al. (2016a) determined Teff = 70 000 ± 2000 K and log g = 7.5 ± 0.1 in a detailed reanalysis of optical and UV spectra. They included 27 elements, namely He, C, N, O, Al, Si, P, S, Ca, Sc, Ti, V, Cr, Mn, Fe, Cr, Ni, Zn, Ga, Ge, As, Kr, Zr, Mo, Sn, Xe, and Ba, in their models. From these, we calculated the EUV spectrum (228 Å ≤ λ ≤ 910 Å) with 1601 atomic levels treated in NLTE, considering 2481 lines of the elements He–S and about 30 million lines of the elements with Z ≥ 20. The frequency grid comprised 174 873 points with Δλ ≤ 0.005 Å.

Figure 2 demonstrates the determination of the interstellar reddening. We apply the reddening data of Morrison & McCammon (1983, provided for 1.26 Åλ ≤ 413 Å and extrapolated toward the He i ground-state threshold) and Fitzpatrick (1999, λ ≥ 911 Å. Between the He i ground-state edge and the H i Lyman edge, only absorption due to H i is considered. To determine EBV, we normalized our models to the 2MASS H brightness (14.766 ± 0.063, Cutri et al. 2003a,b). To match the observed flux level between about 400 Å to 600 Å, EBV = 0.00026 ± 0.00003 is necessary. This is less than EBV = 0.015 ± 0.002 that was used by Rauch et al. (2016a) to reproduce the observed FUSE flux level. With the Galactic reddening law of Groenewegen & Lamers (1989, log (NH i/EBV = 21.58 ± 0.1 and the total cloud column density of interstellar H i of 1.5 ± 0.2 × 1018 cm-2 (measured from L β, Rauch et al. 2016a), we can calculate which is within error limits well in agreement with our result.

A close look at the EUV wavelength range shows still a significant difference between model and observation (Fig. 3, top panel), most prominent between 250 Å and 400 Å and between 504 Å and 550 Å. Our present models reduced the deviation by about a factor of two compared the models of Werner et al. (2001). The EUV problem cannot be solved by using a cooler model, even at Teff = 66 000 K, which is already outside the error range of Teff = 70 000 ± 2000 K given by Rauch et al. (2016a), no sufficient improvement is achieved. The impact of metal opacities is demonstrated in Fig. 3 by a model that considered only opacities from He, C, N, O, Fe, and Ni with same abundance ratios like our best model. To test the impact of additional opacity, we artificially increased the Ni abundance by factor of 200 to match the model’s flux to the observed between 250 Å and 280 Å. This reduced the flux discrepancy between 300 Å and 400 Å as well while the wavelength region above the He i ground-state threshold is unaffected. However, we conclude that even in our advanced models opacity is missing from elements that are hitherto not considered. To include, for example, other trans-iron elements requires detailed laboratory measurements of their spectra and the extensive calculation of transition probabilities.

8. What is the nature of RE 0503–289?

RE 0503–289 was first classified to be a DO-type WD (Barstow et al. 1993). Its optical spectrum exhibits an absorption trough around C IV λλ 4646.62 − 4687.95 Å and He II λ 4685.80 Å. This trough is the spectroscopic criterion for the H-deficient PG 1159-type stars (e.g., Werner & Herwig 2006). Figure 4 shows the comparison of the wavelength region around this trough for the PG 1159 prototype PG 1159−035 (V GW Vir, WD 1159−035, Teff = 140 000 ± 5000 K, log g=7.0 ± 0.5, Jahn et al. 2007) and the O(He)-type WD KPD 0005+5106 (WD 0005+511, Teff = 195 000 ± 15 000 K, log g = 6.7 ± 0.2, Werner & Rauch 2015). Both objects are at an earlier state of stellar evolution than RE 0503–289. The strengths of the PG 1159 absorption troughs are almost the same for the much hotter PG 1159−035 and RE 0503–289, although their photospheric C abundances are significantly different, ≈ 48% by mass (Jahn et al. 2007) and ≈ 2%, respectively. The cool PG 1159-type star PG 0122+200 has about 22% of C in its photosphere (Werner & Rauch 2014).

thumbnail Fig. 4

Section of the optical spectra of PG 1159−035 (from SPY, shifted by 0.3 in flux units), RE 0503–289 (EMMI), and PG 0122+200 (Keck, shifted by −0.3;  from top to bottom) around the PG 1159 absorption trough. For RE 0503–289, the synthetic spectrum of Rauch et al. (2016a) is overplotted (red line). The green, dashed lines indicate the continuum level.

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In a log Tefflog g diagram (Fig. 5), RE 0503–289 is located at the so-called PG 1159 wind limit (Unglaub & Bues 2000, their Fig. 13, digitized with Dexter) that was predicted for a ten-times-reduced mass-loss rate (line A, calculated with = 1.29 × 10-15L1.86 from Bloecker 1995; Pauldrach et al. 1988). This line approximately separates the regions that are populated by PG 1159-type stars and DO-type WDs. Lines B and C in Fig. 5 show where the photospheric C content is reduced by factors of 0.5 and 0.1, respectively, when using the mass-loss rate given above. To the right of line D, no PG 1159 star is located.

Werner et al. (2014) suggested a mass ratio C / He = 0.02 to conserve previously assigned spectroscopic classes. However, PG 1159 stars span a wide range of C / He (0.03 − 0.33, Werner et al. 2014).

RE 0503–289 is located close to line B of Unglaub & Bues (2000) in Fig. 5, that is, its C abundance should be already reduced by a factor of 0.5. Thus, it is likely that RE 0503–289 had a C / He ≈ 0.05 in its antecedent PG 1159-star phase. Even now, its C / He lies a bit higher than 0.02 and RE 0503–289 may be classified as a PG 1159 star as well. This is corroborated by the still high efficiency of radiative levitation that is responsible for the extremely high overabundances of trans-iron elements (Rauch et al. 2016b). However, the transition from a PG 1159-type star to a DO-type star is smooth and RE 0503–289 is an ideal object to study this in detail. Unfortunately, the strong radiative levitation of trans-iron elements wipes out all information about their asymptotic giant branch (AGB) abundances and RE 0503–289 is not suited to constrain AGB nucleosynthesis models.

thumbnail Fig. 5

Location of RE 0503–289 and related objects (Hügelmeyer et al. 2006; Kepler et al. 2016; Reindl et al. 2014b,a; Werner & Herwig 2006) in the log Tefflog g plane (cf., http://www.star.le.ac.uk/~nr152/He.html for stellar parameters). Evolutionary tracks for H-deficient WDs (Althaus et al. 2009) labeled with their respective masses in M are plotted for comparison. Transition limits predicted by Unglaub & Bues (2000) are indicated (see text for details).

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9. Results

RE 0503–289 fulfills criteria of PG 1159 star and of DO-type WD classifications. The presence of the strong PG 1159 absorption trough around He II λ 4685.80 Å (Fig. 4) shows that RE 0503–289 could be classified as a PG 1159 star, although its C abundance would then be the lowest of this group. It is located close to the so-called PG 1159 wind limit (Fig. 5), meaning that it is close to the regime in which gravitation will dominate and pull metals down, out of the atmosphere. The strongly increased abundances of trans-iron elements, however, indicate that radiative levitation is still efficiently counteracting this process. Thus, RE 0503–289 has not arrived in its final stage of evolution. Formally, due to its log g> 7, the DO-type WD classification is right.

In the observed spectra, we identified 1272 lines in the wavelength range from the extreme ultraviolet to the near infrared. 287 lines remain unidentified. The best model of Rauch et al. (2016a) reproduces well most of the identified lines.

The EUV problem (Sect. 7) – the difference between observed and synthetic flux in the EUV is still present. Our advanced model atmospheres include opacities of 27 metals but their flux in the EUV is still partly about a factor of approximately two too high compared with the observation. We expect that missing metal opacities are the reason for this discrepancy.


6

IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Associated Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

Acknowledgments

D.H. and T.R. are supported by the German Aerospace Center (DLR, grants 50 OR 1501 and 05 OR 1507, respectively). The German Astrophysical Virtual Observatory (GAVO) project at Tübingen had been supported by the Federal Ministry of Education and Research (BMBF, 05 AC 6 VTB, 05 AC 11 VTB). Financial support from the Belgian FRS-FNRS is also acknowledged. P.Q. is research director of this organization. Some of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NNX09AF08G and by other grants and contracts. We thank Ralf Napiwotzki for putting the reduced ESO/VLT spectra at our disposal. The TEUV tool (http://astro-uni-tuebingen.de/~TEUV), the TGRED tool (http://astro-uni-tuebingen.de/~TGRED), the TIRO service (http://astro-uni-tuebingen.de/~TIRO), the TMAD service (http://astro-uni-tuebingen.de/~TMAD), the TOSS service (http://astro-uni-tuebingen.de/~TOSS), and the TVIS tool (http://astro-uni-tuebingen.de/~TVIS) used for this paper were constructed as part of the activities of the German Astrophysical Virtual Observatory. This work used the profile-fitting procedure OWENS developed by M. Lemoine and the FUSE French Team. This research has made use of NASA’s Astrophysics Data System and of the SIMBAD database operated at CDS, Strasbourg, France.

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Appendix A: Identified and unidentified lines in the spectrum of RE 0503–289

Tables A.1–A.5 are available at the CDS.

Appendix B: Observed spectra of RE 0503–289 compared with our best model

In the following figures, we show the comparison of our synthetic spectra with the FUSE (Fig. B.1, HST/STIS (Fig. B.2, and optical (Fig. B.3 observations. A visualization via TVIS is available at http://astro.uni-tuebingen.de/~TVIS/objects/RE0503-289

thumbnail Fig. B.1

FUSE observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER

thumbnail Fig. B.2

HST/STIS observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER

thumbnail Fig. B.3

Optical (SPY) observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER

Appendix C: WWW interfaces of TEUV, TGRED, and TVIS

thumbnail Fig. C.1

TEUV WWW interface.

Open with DEXTER

thumbnail Fig. C.2

TGRED WWW interface.

Open with DEXTER

thumbnail Fig. C.3

TVIS WWW interface.

Open with DEXTER

All Tables

Table 1

History of Teff and log g determinations (cf., Müller-Ringat 2013).

Table 2

Photospheric abundances (mass fraction) of RE 0503–289.

Table 3

Newly calculated transition probabilities.

Table 4

Statistics of the identified (in brackets: newly identified in this paper) and unidentified lines in the observed spectra.

All Figures

thumbnail Fig. 1

Determination of vrad from individual lines in the FUSE (top panel) and HST/STIS observations (bottom). The full horizontal lines indicate the average vrad for FUSE and HST/STIS, respectively. The dashed lines show the 1σ error.

Open with DEXTER
In the text
thumbnail Fig. 2

Determination of EBV. Top: reddening with EBV = 0.00026 applied to our synthetic spectrum in the wavelength range from the EUV to the NIR. Bottom: same like top panel, for EBV = 0.00016 (blue), 0.00026 (red), and 0.00036 (green) in the EUV wavelength range. Prominent lines are marked.

Open with DEXTER
In the text
thumbnail Fig. 3

Comparison of the EUVE observation (gray line in both panels) with our models. Top panel: two models with Teff = 70 000 K (red) and Teff = 66 000 K (blue). Identified photospheric lines are marked at the top. Bottom panel: three models with Teff = 70 000 K. Red, thick line: model from the top panel, red, dashed line: model with 200 times increased Ni abundance, blue, thin line: model that considered only opacities of He, C, N, O, Fe, and Ni.

Open with DEXTER
In the text
thumbnail Fig. 4

Section of the optical spectra of PG 1159−035 (from SPY, shifted by 0.3 in flux units), RE 0503–289 (EMMI), and PG 0122+200 (Keck, shifted by −0.3;  from top to bottom) around the PG 1159 absorption trough. For RE 0503–289, the synthetic spectrum of Rauch et al. (2016a) is overplotted (red line). The green, dashed lines indicate the continuum level.

Open with DEXTER
In the text
thumbnail Fig. 5

Location of RE 0503–289 and related objects (Hügelmeyer et al. 2006; Kepler et al. 2016; Reindl et al. 2014b,a; Werner & Herwig 2006) in the log Tefflog g plane (cf., http://www.star.le.ac.uk/~nr152/He.html for stellar parameters). Evolutionary tracks for H-deficient WDs (Althaus et al. 2009) labeled with their respective masses in M are plotted for comparison. Transition limits predicted by Unglaub & Bues (2000) are indicated (see text for details).

Open with DEXTER
In the text
thumbnail Fig. B.1

FUSE observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER
In the text
thumbnail Fig. B.2

HST/STIS observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER
In the text
thumbnail Fig. B.3

Optical (SPY) observation (gray) compared with the best model (red). Stellar lines are identified at top. “unid.” denotes unidentified lines.

Open with DEXTER
In the text
thumbnail Fig. C.2

TGRED WWW interface.

Open with DEXTER
In the text

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