A&A 616, A73 (2018) Astronomy https://doi.org/ 10.1051/0004-6361/201833200 & © ESO 2018 Astrophysics

The hot in the peculiar binary nucleus of the planetary EGB 6? K. Werner1, T. Rauch1, and J. W. Kruk2

1 Institute for Astronomy and Astrophysics, Kepler Center for Astro and Particle Physics, Eberhard Karls University, Sand 1, 72076 Tübingen, Germany e-mail: [email protected] 2 NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA

Received 10 April 2018 / Accepted 14 May 2018

ABSTRACT

EGB 6 is an extended, faint old (PN) with an enigmatic nucleus. The central (PG 0950+139) is a hot DAOZ-type white dwarf (WD). An unresolved, compact emission knot was discovered to be located 000. 166 away from the WD and it was shown to be centered around a dust-enshrouded low-luminosity star. It was argued that the dust disk and evaporated gas (photoionized by the hot WD) around the companion are remnants of a disk formed by wind material captured from the WD progenitor when it was an asymptotic giant branch (AGB) star. In this paper, we assess the hot WD to determine its atmospheric and stellar parameters. We performed a model-atmosphere analysis of ultraviolet (UV) and optical spectra. We found Teff = 105 000 5000 K, log g = 7.4 0.4, and a solar helium abundance (He = 0.25 0.1, mass fraction). We measured the abundances of ten more± species (C, N, O,± F, Si, P, S, Ar, Fe, Ni) and found essentially solar± abundance values, indicating that radiation-driven wind mass-loss, with a theoretical ˙ +1.1 rate of log(M/M /yr) = 11.0 0.8, prevents the gravitational separation of elements in the photosphere. The WD has a mass of +0.12 − − +1.26 M/M = 0.58 0.04 and its post-AGB age (log(tevol/yr = 3.60 0.09)) is compatible with the PN kinematical age of log(tPN/yr = 4.2). In addi- tion, we examined− the UV spectrum of the hot nucleus of− a similar object with a compact emission region, Tol 26 (PN G298.0+34.8), and found that it is a slightly cooler DAOZ WD (Teff 85 000 K), but this WD shows signatures of gravitational settling of heavy elements. ≈ Key words. planetary nebulae: general – : abundances – stars: atmospheres – stars: AGB and post-AGB – white dwarfs

1. Introduction WD progenitor’s wind during its asymptotic giant branch (AGB) stage, and that this disk is surrounded by gas (the CEK) pho- EGB 6 (PN G221.5+46.3) is a faint old planetary nebula (PN) toionized by the hot WD. The nature of the companion is in fact with a mysterious nucleus. The PN was discovered by Ellis et al. unknown because it is enshrouded by dust (Bond et al. 2016). It (1984) and has an angular extension of 110 130. At a nomi- must be a low-luminosity object, a dM star or even another WD. nal distance of 725 pc, the linear diameter of× the PN is 1 pc. ≈ ≈ For a detailed summary of our current knowledge about this sys- The central star (PG 0950+139) is a very hot white dwarf (WD) tem, we refer to the exhaustive paper by Bond et al.(2016) and that was independently discovered by Ellis et al.(1984) and the references therein. Palomar Green (PG) survey (Green et al. 1986). Bright nebula In the present paper, we investigate in detail the hot cen- emission lines superposed on the WD spectra were shown not tral star of EGB 6. It was classified as a hot DA white dwarf to stem from the old, low-surface brightness PN, but from a by Fleming et al.(1986). Later on, by the detection of ionized relatively dense (n = 2.2 106 cm 3; Dopita & Liebert 1989), e × − helium, Liebert et al.(1989) established the star as a DAO white unresolved compact emission knot (CEK) that was thought to dwarf, and after the detection of metal lines in the ultraviolet be centered on the WD. An infrared excess was discovered by (Gianninas et al. 2010) the star was assigned to the DAOZ class. Zuckerman et al.(1991) and attributed to the presence of an The analysis of the optical spectrum of the WD is M-type stellar companion. It came as a big surprise when imag- hampered by strong emission lines from the CEK. Using pure- ing and grism spectroscopy obtained with the Hubble Space hydrogen local thermodynamic equilibrium (LTE) model atmo- Telescope (HST) revealed that the CEK is a point-like source spheres, Liebert et al.(1989) derived Teff = 70 000 7000 K and at the location of the companion, 000. 166 away from the WD log g = 7.5 0.25 from a fit to the Hγ line only. From± new obser- (Liebert et al. 2013), corresponding to a projected linear sep- vations, and± with pure hydrogen non-LTE models for Balmer aration of 118 AU. The origin of the material making up the ≈ line fits, Liebert et al.(2005) found Teff = 108 390 16 687 K CEK is unknown. It was suggested that it is the remnant of a and log g = 7.39 0.38. ± disk containing material captured by the companion star from the Using non-LTE± model atmospheres, Gianninas et al.(2010) derived T = 93 230 20 425 K and log g = 7.36 0.79, empha- ? Based on observations made with the NASA-CNES-CSA Far eff ± ± Ultraviolet Spectroscopic Explorer and the NASA/ESA Hubble Space sizing increased uncertainties in the measured parameters due Telescope, obtained at the Space Telescope Science Institute, which to the exclusion of the Balmer line cores from the fitting pro- is operated by the Association of Universities for Research in cedure. For the first time, they measured the helium abundance +0.7 Astronomy, Inc., under NASA contract NAS5-26666. (He/H = 0.1 0.9, by mass) which, at face value, is lower than the − Article published by EDP Sciences A73, page 1 of9 A&A 616, A73 (2018) solar abundance (He/H = 0.34), but highly uncertain. In their 2002; Hébrard & Moos 2003). It has the capacity to consider hydrogen plus helium models, in order to compute a more realis- several, individual ISM clouds with their own radial and turbu- tic atmospheric structure, they included the elements C, N, and O lent velocity, temperature, chemical composition, and respective in solar amounts, as prescribed by the analysis of Werner(1996), column densities. We identified and modeled lines of H I,D I, who outlined the solution to the Balmer line problem. This mea- H2 (J = 0–4), HD (J = 0), C II-III,C II∗,N I–III,O I, Si II,P II, sure was motivated by the fact that they identified metal lines in S III, Ar I, and Fe II. Far Ultraviolet Spectroscopic Explorer (FUSE) spectra; however, The star was observed in 1993 with the Faint Object Spec- abundance determinations were not performed. trograph (FOS) aboard HST. Two spectra were taken with Here, we present a new spectral analysis of the WD. Our gratings G160L (1150–2500 Å, resolution 7 Å) and G270H ≈ primary aim was to determine the metal abundances from the (2220–3300 Å, resolution 2 Å). Because of the low resolution ultraviolet (UV) spectra taken with FUSE and to improve the of the G160L spectrum, even≈ the strongest spectral lines pre- measurement of Teff, log g, and He abundance employing, in dicted by our models (He II λ1640, C IV λ1550, N V λ1240, addition, previously unused optical spectra. It was expected that O V λ1371) are not detected beyond doubt. In the G270H spec- the results would contribute one piece to resolve the puzzle of trum, three members of the He II Fowler series can be seen 1 this extraordinary PN nucleus . (n = 3 5, 6, 7). It has been claimed that EGB 6 is the prototype of a small Three→ observations during 1987–1989 were performed with class of PNe with central compact emission regions (Frew & IUE (SWP camera, resolution 5 Å). The co-added spectrum was Parker 2010). Related to them could be Tol 26 (PN G298.0+34.8; taken from the database created by Holberg et al.(2003). Its flux Hawley 1981), whose central star was hitherto unexplored. From level matches the HST/FOS data. At best, He II λ1640 and C IV the nebula analysis, Teff > 50 000 K was concluded. Here we λ1550 are barely detectable. present the investigation of an archival UV spectrum obtained The UV spectra up to λ = 1770 Å are shown in Fig.2. We with the Cosmic Origins Spectrograph (COS) aboard HST, to fit the continuum shape with our final photospheric model check for the possible similarity between the hot central stars of assuming interstellar reddening E(B V) = 0.025 0.005, EGB 6 and Tol 26. in good agreement with the value derived− by Bond± et al. (2016; E(B V) = 0.02) and confirming the conclusion that − 2. Observations the hot central star suffers no reddening within the system because the total reddening in the direction of EGB 6 is about For our analysis of EGB 6, we use UV spectra which were E(B V) = 0.027 (Schlafly & Finkbeiner 2011). From the Lyman observed with the International Ultraviolet Explorer (IUE), lines,− we derive an interstellar neutral hydrogen column density 2 HST, and FUSE. Optical spectra were obtained by the Sloan of log(nH/cm ) = 20.3 0.4. This mainly follows from fitting the Digital Sky Survey (SDSS) and the Supernovae Type Ia Lyα profile because only± for this line is the equivalent width Progenitor Survey (SPY; Napiwotzki et al. 2003). For Tol 26, we dominated by the interstellar medium (ISM) and not by the used the HST/COS spectrum. photospheric contribution. A HST/COS spectrum of the central star of Tol 26 was 2.1. Ultraviolet spectra recorded in 2014 with grating G140L. It has a useful wavelength range of 1160–2000 Å and a spectral resolution of 0.6 Å. It EGB 6 was observed by FUSE twice in 2000 and we retrieved is discussed≈ in detail in Sect. 3.4. The spectrum was≈ smoothed the spectra from the MAST archive. The first observation with a 0.1 Å wide boxcar and the models were folded with 0.6 Å (A0340101) had imprecise target coordinates, so the star was on Gaussians. the edge of the slit in two channels and completely missed the other two. The second observation (A0341101) was good in all channels, so we scaled the usable data from the first observation 2.2. Optical spectra to match the second. The reduction procedure was the same as The archival SDSS spectrum of the EGB 6 nucleus covers the that described in Werner et al.(2016). The resulting spectrum wavelength range 3600–10 300 Å with a resolution of 2.5 Å. is shown in Fig.1. The resolving power is about 20 000. We The spectrum is a co-addition of four single observations.≈ We smoothed the spectrum by a 0.05 Å wide boxcar and, accord- rectified it for our fitting with normalized synthetic line profiles ingly, our synthetic spectra were convolved with a Gaussian with (Fig.3). The only photospheric lines that can be identified are FWHM = 0.07 Å. from H and He II. The cores of the Balmer lines and He II λ4686 The hydrogen Lyman lines are blended with the respective (and perhaps also He II λ5412) are contaminated with emission He II line series. From the latter, three isolated lines are visible, lines from the CEK around the companion (Liebert et al. 2013), with decreasing strength, at 1085, 992, and 959 Å (transitions hampering the determination of atmospheric parameters. between levels with principal quantum numbers n = 2 5, 6, 7). → To mitigate this problem, we used SPY spectra, which have a The FUSE spectrum was published by Gianninas et al.(2010) higher resolution ( 0.1 Å), allowing us to trace the photospheric and photospheric lines of seven metals were identified (C, N, O, line profiles closer≈ toward the line core, because the nebular Si, P, S, Fe); however, it was not subject to a quantitative analysis. emission lines are rather narrow. However, the signal-to-noise We identified many more lines of these elements plus three more ratio (S/N) of these spectra is relatively poor. In total, three SPY species (F, Ar, and Ni; see Sect. 3.2). spectra (each consisting of three segments covering the ranges The FUSE observations exhibit many absorption lines of 3400–4520 Å, 4620–5600 Å, and 5680–6640 Å) were used and the ISM. To unambiguously identify stellar lines, we employed rectified. We abandoned co-addition because of the very differ- the program OWENS (Lemoine et al. 2002; Hébrard et al. ent S/N of the single spectra in different wavelength regions. In 1 To pick up the thread from Bond et al.(2016), a fitting quotation of the SPY spectra, too, no metal lines are detectable. the fictional Sherlock Holmes might be: “It is, of course, a trifle, but The He II λ3203 line, covered by the HST/FOS G270H there is nothing so important as trifles” (Conan Doyle 1891). spectrum, is obviously not affected by nebula emission. The

A73, page 2 of9 K. Werner et al.: The hot white dwarfA&A in proofs: the peculiarmanuscript binary no. nucleus aa of the planetary nebula EGB 6

EGB 6 O IV N IV He II S VI N IV 1 Ly η Ly ζ Ly ε Ly δ Ly γ

O V O V 0 920 930 940 950 960 970 980

O VI O IV Fe VI O V He II O V He II O VI 1 N V Ly β

S VI O V O IV O V 0 990 1000 1010 1020 1030 1040

relative flux Fe VII FeVIII Si IV 1 C IV He II

O V O IV Ni VI S VI O IV O IV ArVII O V P V 0 O VI 1060 1070 1080 1090 1100 1110 P V Fe VII Si IV 1 Fe VI C IV

Fe VIII F VI Fe VIII FeVI Ni VI O VI O IV

0 1130 1140 1150 1160 1170 1180 o wavelength / A

Fig. 1. 1.FUSEFUSE spectrum spectrum of of the the white white dwarf dwarf in in EGB EGB 6 6 (gray) (gray) compared compared to to a photospherica photospheric model model spectrum spectrum (red; (red;TeffTe=ff 105= 105 000 000 K, K,log logg =g 7.4)= 7.4) with with the adoptedthe adopted parameters parameters (see (see Table Table1). The 1). The same same model model attenuated attenuated by interstellar by interstellar lines lines is plotted is plotted in blue. in blue. Prominent Prominent photospheric photospheric lines lines are are identified. identified. evaluatedsolar abundance optical (He and/H near-UV=0.34), but H and highly He uncertain.II lines are In shown their hy- in theperformed results would line formation contribute iterations one piece (i.e., to resolve keeping the the puzzle atmo-of Fig.drogen3. plus helium models, in order to compute a more realistic thisspheric extraordinary structure fixed) PN nucleus. for argon1 and fluorine using the model atmospheric structure, they included the elements C, N, and O in atomsIt has presented been claimed in Werner that et EGB6 al.(2015 is). the prototype of a small We first turn to the WD in EGB 6 (Sects. 3.1–3.3) and then 3.solar Spectral amounts, analysis as prescribed by the analysis of Werner (1996), class of PNe with central compact emission regions (Frew & who outlined the solution to the Balmer line problem. This mea- Parkerto the nucleus 2010).Related in Tol 26 to (Sect. them could3.4). be Tol26 (PNG298.0+34.8, Wesure used was themotivated Tübingen by the Model-Atmosphere fact that they identified Package metal (TMAP line2)s Hawley 1981), whose central star was hitherto unexplored. From toin computeFar Ultraviolet non-LTE, Spectroscopic plane-parallel, Explorer line-blanketed (FUSE) spectra; atmospherehow- the3.1. nebulaEffective analysis, temperature,Teff > surface50000K gravity, was concluded. He abundance Here we modelsever, abundance in radiative determinations and hydrostatic were not equilibrium performed. (Werner & present the investigation of an archival UV spectrum obtained We started our spectrum fitting procedure for EGB 6 from DreizlerHere, 1999 we; present Werner a et new al. spectral2003, 2012 analysis). The of models the WD. include Our with the Cosmic Origins Spectrograph (COS) aboard HST, to a model with temperature, gravity, and helium abundance H,primary He, C, aim N, O, was Si, to P, determine S, Fe, and theNi. metalThe employed abundances model from atoms the close to the values derived by Gianninas et al.(2010): Teff = areultraviolet described (UV) in detail spectra by taken Werner with etFUSE al.(2018 and). In to addition, improve we the 1 93To 000 pick K, uplog theg thread= 7.4, fromH = 0.9,Bond and et al. He (2016), = 0.09 a fitting (mass quotati fractions),on of measurement of T , log g , and He abundance employing, in 2 http://astro.uni-tuebingen.de/~TMAPeff theand fictional assuming Sherlock solar Holmes abundances might for be: the “It is, metals. of course, Keeping a trifle, fixed but addition, previously unused optical spectra. It was expected that there is nothing so important as trifles.” (Conan Doyle 1891). A73, page 3 of9 Article number, page 2 of 9 K. Werner et al.: The hot white dwarfA&A in the 616,peculiar A73 binary (2018) nucleus of the planetary nebula EGB 6

He II 8 H I Lyman -1 o A

-1 6 EGB 6 s -2 20 -2 Teff =105 000 K log g=7.4 E(B-V)=0.025 nH =2.0x10 cm

erg cm 4 -13 / 10 λ f 2 FUSEHST

1000 1100 1200 1300 1400 1500 1600 1700 o wavelength / A

Fig. 2. FUSEFUSE spectrum spectrum (gray) (gray) and and HST HST/FOS/FOS spectrum spectrum (red) (red) of ofthe the white white dwarf dwarf in EGB in EGB 6 compared 6 compared to a photo atospheric photospheric model model spectrum spectrum (blue; (blue;Teff = g . T105eff = 000 105 K, 000 log K, log= 7.4)g = with 7.4) withthe adopted the adopted parameters parameters (see(see Table Table 1). A1). reddening A reddening of E of(BE(BV)V=) =0 0025.025 waswas applied applied to to the the model model to to fit fit the the observed − − 20 2 continuumcontinuum shape. The model is attenuated by interstellar LymanLyman line absorption (column density nH = 22 10 cm− ). The black dashed line is the H × unattenuated model. For clarity, the FUSE spectrum and the mmodelsodels in the respective wavelength range were smoothed with a Gaussian with 1 Å FWHMFWHM.. The The model model in in the the HST/FOS HST/FOS range range was was convolved convolved with with a a 7 7 Å Å Gaussian. Gaussian.

logcheckg and for the the possible H/He ratio, similarity we then between adjusted the iteratively hot centralT steffarsand of individualstrengths. The ISM central clouds emission with their in own He radialII λ4686 and is turbulent of nebular veloc- ori- theEGB6 metal and abundances Tol26. to achieve a good fit to the metal lines ity,gin temperature,because its height chemical is not composition, reached by the and emission respective reversal column of in the FUSE spectrum. From the relative strengths of lines densities.the photospheric We identified line profile. and modeled lines of H i, D i, H2 (J = from different ionization stages of nitrogen, oxygen, and iron 0–4),Comparing HD (J = 0), our C ii results–iii, C ii with∗, N i– theiii, O mosti, Si ii recent, P ii, S onesiii, Ar byi, (see2. Observations Sect. 3.2), we realized that the effective temperature must andGianninas Fe ii. et al.(2010; see Sect.1), the main progress we be higher, namely 105 000 K, within a narrow margin of about achieved is the rather tight constraint of the effective temper- For our analysis of EGB6, we use UV spectra which were ob- The star was observed in 1993 with the Faint Object Spec- 5000 K. Then the possible range of the surface gravity was ature ( 5% compared to 22%) by using ionization balances served with the International Ultraviolet Explorer (IUE), HST, trograph (FOS) aboard HST. Two spectra were taken with grat- checked with the optical H and He lines, confirming the initial of metals.± In addition, the± error in the surface gravity could be and FUSE. Optical spectra were obtained by the Sloan Digi- ings G160L (1150–2500Å, resolution 7Å) and G270H(2220– valuetal Sky within Surveylog (SDSS)g = 7.4 and0. the4. The Supernovae UV lines Type of H Ia and Progenitor He are reduced from about 0.8 to 0.4 dex.≈ The same holds for the compatible with this result± but they are less sensitive to log g. 3300Å,He/H abundance resolution ratio.±2Å). As a Because consequence,± of the stellar low resolution parameters of can the Survey (SPY, Napiwotzki et al. 2003). For Tol26, we used the G160L spectrum, even≈ the strongest spectral lines predicted by ThenHST/COS the He spectrum. abundance was refined by using the optical and UV be determined with higher precision (Sect. 3.3). spectra and we found a better fit with a higher He/H ratio, namely our models (He ii λ1640, C iv λ1550, N v λ1240, O v λ1371) are not detected beyond doubt. In the G270H spectrum, three mem- a solar value (H = 0.74 and He = 0.25). A model computed with 3.2. Metal abundances bers of the He ii Fowler series can be seen (n = 3 5, 6, 7). log2.1.g Ultraviolet= 8.0 shows spectra H and He lines that are significantly too broad. → To investigate the effects of the uncertainty in log g on WeThree detected observations lines from during ten 1987–1989 metals in werethe FUSE performed spectrum with EGB6 was observed by FUSE twice in 2000 and we re- (Figs.1 and4). These lines are well known from previous anal- the other atmospheric parameters, we repeated the fitting proce- IUE (SWP camera, resolution 5Å). The co-added spectrum was trieved the spectra from the MAST archive. The first observation yses. Line lists, from which detailed laboratory wavelengths and dure with log g = 7.0 models. The lower gravity leads to lower taken from the database created by Holberg et al. (2003). Its flux (A0340101) had imprecise target coordinates, so the star was on transitions can be obtained, were presented by Werner et al. particle densities and, according to the Saha equation, metal ion- level matches the HST/FOS data. At best, He ii λ1640 and C iv the edge of the slit in two channels and completely missed the (2015, their Tables 2–5; 2017, their Tables 2 and 3). The derived izationother two. balances The second are shifted observation to higher (A0341101) stages. To compensatewas good in for all λ1550 are barely detectable. this, T must be reduced to retain good line fits, namely down to abundances are listed in Table1 and displayed in Fig.5. channels,eff so we scaled the usable data from the first observa- TheFrom UV carbon, spectra two up C toIVλdoublets= 1770Å are are seen shown in the in FUSE Fig.2. spec- We 100tion 000 to match K. Metal the second.abundances The arereduction hardly procedure affected by was this the change same T g fittrum, the namelycontinuum the shape 3p–4d withand our 3d final–4 f transitions photospheric at model 1108 and as- inas thateff and describedlog , in but Werner the helium et al. abundance (2016). The must resulting be reduced spec- in the model. Otherwise, the He II lines become too strong, suming1169 Å, respectively.interstellar reddening In our finalE( modelB V with) = C0. =025 0.003,0. the005, line in trum is shown in Fig.1. The resolving power is about 20000. goodcores agreementare not deep with enough the value to match derived− the observation. by Bond et± A al. detailed (2016, inWe particular smoothed those the spectrum at 959, 4686, by a and 0.05Å 5412 wide Å, and boxcar a detectable and, ac- λ Einspection(B V) = shows0.02) and that confirming the computed the conclusion line profiles that show the hot central cen- cordingly,4542 line our is predicted synthetic that spectra is not were seen convolved in the observations. with a Gaussian We − arrive at He = 0.18 (and thus H = 0.81). From these experiments tralemission star su cores.ffers no This reddening is unrealistic, within the but system despite because some the numeri- total with FWHM=0.07Å.= g . . reddeningcal testing in concerning the direction the of temperature EGB6 is about structureE(B ofV the) = model,0.027 we conclude Teff 105 000 5000 K, log = 7 4 0 4, and − He =The0.25 hydrogen0.1. The Lyman error lines in± the are individual blended with metal the abundances± respective (Schlaflythe reason & remains Finkbeiner unknown. 2011). IncreasingFrom the Lyman the carbon lines, abundance we derive He ii line series.± From the latter, three isolated lines are visible, 2 is estimated not to exceed 0.5 dex. anmitigates interstellar the problemneutral hydrogencolumn in the line cores, density but leads of log( ton tooH/cm strong)= withIn decreasing Fig.3, we comparestrength,± the at 1085Å, observed 992Å, optical and and 959Å near-UV (transi- H 20line.3 wings0.4. Thisand to mainly the appearance follows from of fitting optical the C IV Lyαlines,profile which be- tions between levels with principal quantum numbers n = 2 ± and He lines with our final model. Although nebular emission causeare not only observed. for this lineThe is variations the equivalent in log widthg, Te dominatedff, and He by abun-the , , → lines5 6 7). mask The the FUSE Balmer spectrum line centers, was published we note a by Balmer Gianninas line prob- et al. interstellardance described medium in Sect. (ISM) 3.1 and do not improveby the photospheric the C IV line contr fits.ibu- The lem(2010) (Napiwotzki and photospheric 1993), that lines is, of the seven depth metals of innermost were identifie partsd (C, of tion.same problem was encountered in our analysis of a hot DO WD N, O, Si, P, S, Fe); however, it was not subject to a quantitative with similar temperature and gravity (PG 1034+001, 115 000/7.0; the observed photospheric profiles of Hα and Hβ is not quite A HST/COS spectrum of the central star of Tol26 was matchedanalysis. by We the identified models. many Gianninas more etlines al.( of2010 these) were elements unable plu tos Werner et al. 2017, their Fig. 1). The star is too hot to show the three more species (F, Ar, and Ni; see Sect.3.2). recorded in 2014 with grating G140L. It has a useful wavelength detect this problem because of the lower resolution of their spec- C III multiplet at 1175 Å, which is often seen in other DAO WDs The FUSE observations exhibit many absorption lines of the range of 1160–2000Å and a spectral resolution of 0.6Å. It tra. In order to fit these lines, a lower temperature would be and in Tol≈ 26 (see Sect. 3.4). ≈ ISM. To unambiguously identify stellar lines, we employed the is discussed in detail in Sect.3.4. The spectrum was smoothed required; however, this would give too-deep Hγ and Hδ line cores Nitrogen displays the N IV multiplet at 922–924 Å and the program OWENS (Lemoine et al. 2002; Hébrard et al. 2002; with a 0.1Å wide boxcar and the models were folded with 0.6Å and, as mentioned above, would be in conflict with the metal line singlet at 955 Å. We also see two N V doublets at 1048 and Hébrard & Moos 2003). It has the capacity to consider several, Gaussians. A73, page 4 of9 Article number, page 3 of 9 K. Werner et al.: The hot white dwarfA&A in proofs: the peculiarmanuscript binary no. nucleus aa of the planetary nebula EGB 6

componentThe He ii isλ blended3203 line, by covered one of bythe the mentioned HST/FOS relatively G270H broadspec- α H trum,photospheric is obviously O VI lines. not aff Anotherected by strong nebula Si emission.IV doublet The at 1067 evalu- Å atedis blended optical by and a near-UV strong interstellar H and He ii line.lines The are phosphorus shown in Fig. abun- 3. dance was derived from the P V λλ1118/1128 resonance doublet. λλ Hβ The sulfur abundance was found from the S VI 933/945 res- 3.onance Spectral doublet, analysis and from two weaker subordinate S VI lines at 1000 and 1118 Å. The argon abundance was measured from the We used the Tübingen Model-Atmosphere Package (TMAP2) 1.0 rather prominent Ar VII λ1063.55 line, which is blended by the Hγ towing compute of an adjacent non-LTE, interstellar plane-parallel, line. line-blanketed atmosphere models in radiative and hydrostatic equilibrium (Werner & Drei- The strongest iron lines stem from Fe VII, followed by zler 1999; Werner et al. 2003, 2012). The models include H, He, Fe VI and Fe VIII lines (the only unblended Fe VIII line is C, N, O, Si, P, S, Fe, and Ni. The employed model atoms are at 1148 Å). The relative strengths of lines from the differ- Hδ described in detail by Werner et al. (2018). In addition, we per- formedent ionization line formation stages are iterations sensitive (i.e., temperature keeping the indicators. atmospheric For g structureexample, fixed) at Teff for= 95 argon 000 K and (and fluorinelog using= 7.0), the the model Fe VII atomslines presentedbecome too in Wernerweak, while et al. the(2015). Fe VI lines become too strong. At T = 115 000 K (and log g = 7.4), the Fe VI lines are too weak, 0.5 eff We first turn to the WD in EGB6 (Sects.3.1–3.3) and then HeII 4686 towhile the nucleus the Fe VIII in Tol26lines (Sect.3.4).are too strong. The nickel abundance is

relative flux derived from a multitude of relatively weak Ni VI lines. Other iron group elements were not detected, and we can thus 3.1.exclude Effective extreme temperature, overabundances surface that gravity, were seen He abundance in other DAOs, HeII 5412 namely HS 2115+1148 (Werner et al. 2018) and BD 22◦3467, Wethe excitingstarted our star spectrum of Abell 35 fitting (Ziegler procedure et al. 2012 for). EGB6 Likewise,− from we a modelfound nowith hint temperature, of trans-iron gravity, group and elements helium that abundance are present clos ine = tostrongly the values oversolar derived amounts by Gianninas in several et al. hot (2010): DO WDsTeff (Hoyer93000K, et al. HeII 3203 g = = = log2018, and7.4, references H 0.9 and therein). He 0.09 (mass fractions), and assum- 0.0 g ingTo solar summarize, abundances the for abundances the metals. of Keeping hydrogen, fixed helium, log and and / theten H moreHe identifiedratio, wethen species adjusted in the iteratively white dwarfTeff are,and within the metal error abundances to achieve a good fit to the metal lines in the FUSE HeII 2734 limits, compatible with solar values. spectrum. From the relative strengths of lines from different ion- ization stages of nitrogen, oxygen, and iron (see Sect.3.2), we realized3.3. Stellar that parameters, the effective distance temperature must be higher, namely EGB 6 105000K, within a narrow margin of about 5000K. Then the Stellar parameters are estimated from the comparison of the -30 -15 0 15 30 possible range of the surface gravity was checked with the op- o determined atmospheric parameters Teff and log g with evo- ∆λ / A tical H and He lines, confirming the initial value within log g = 7lutionary.4 0.4. The tracks UV by lines Miller of Hand Bertolami He are(2016 compatible; Fig.7). with We this find re- M/M = . +0.12 L/L = +300 Fig. 3. Observed Balmer and He iiII lineslines of of the the white white dwarf in EGB 6 sult± but they0 58 are0.04 lessand sensitive to72 log48g .. Then The the errors He abundance are domi- − − g (red: SDSS spectrum, gray: SPY spectrum; two lowest lines: HHSTST spec- wasnated refined by the by using uncertainty the optical in log and UV. The spectra post-AGB and we age found is trum) compared to a model spectrum (blue graph; T = 105 000 K, trum) compared to a model spectrum (blue graph; Teeffff = 105 000 K, aabout better 4000 fit with yr; ahowever, higher He the/H mass ratio, uncertainty namely a solar causes value a (H large= log g = 7.4) with the adopted parameters (see Table1 1).). Nebular Nebular emis- emis- t . +1.26 0.74possible and He error:= 0.25).log( Aevol model/yr) = computed3 60 0.09 with. The log mass-lossg = 8.0 shows rates sion lines lines in in the the photospheric photospheric Balmer Balmer line cores line cores are truncat are truncateded for clarity. for − Halong and Hethe lines evolution that are assumed significantly by Miller toobroad. Bertolami(2016) imply clarity. ˙ / . +1.1 log(M M /yr) = 11 0 0.8. g To investigate − the e−ffects of the uncertainty in log on the otherThe atmospheric spectroscopic parameters, distance wed repeatedwas found the by fitting comparing procedure the dereddened visual magnitude V with the respective model 1050 Å, which are partially blended by interstellar lines. The rel- with log g = 7.0 models. The lower0 gravity leads to lower par- 2.2. Optical spectra atmosphere flux, resulting in the relation ative strength of the N IV and N V lines is well matched by our ticle densities and, according to the Saha equation, metal ion- final model. Values of Teff as high as 115 000 K (at log g = 7.4) ization balances are shifted to higher stages. To compensate for The archival SDSS spectrum of the EGB6 nucleus covers the 4 0.4V0 log g or as low as 95 000 K (at log g = 7.0) can be strictly excluded this,d[pc]Te=ff 7must.11 be10 reducedHν toM retain10 good− line, fits, namely down to × × × wavelengthbecause the rangeN IV lines 3600–10300Å are much too with weak a resolution or too strong, of respec-2.5Å. 100000K. Metal abundances are hardly affected by this change ≈ p 3 2 1 1 Thetively, spectrum while the is N aV co-additionline strengths of four are almost single unaffectedobservations. byW thee inwhereTeff Handν = log1.60g ,10 but− theerg helium cm− s− abundanceHz− is the must Eddington be reduced flux rectified it for our fitting with normalized synthetic line profiles × Teff variations. inof the the model. model Otherwise, at 5400 Å, the and HeMii linesis the become stellar too mass strong, in M in. (Fig.3).Prominent The only oxygen photospheric lines from lines three that ionization can be identified stages aarere particularFor our WD, those we at 959Å, have V 4686Å,= 15.998 and (Bond 5412Å, et and al. 2016 a detectable). From ii ii fromdetected. H and The He O.IV The/O coresV line of strengththe Balmer ratio lines is and very He sensitiveλ4686 λE4542(B lineV) = is0. predicted025, we derivedthat is not the seen visual in the extinction observations. using W thee (and perhaps also He ii λ5412) are contaminated with emission − against Teff variations. While the O V lines hardly change within arrivestandard at He relation= 0.18AV (and= 3 thus.1E( HB = V0.81).) = 0 From.077, thesehence experimentsV0 = 15.92. +1100 − linesa range from of the 95 000–110CEK around 000 K, the the companion O IV lines (Liebert quickly et disappearal. 2013), weWe conclude found d =Te660ff = 105 000pc. This5000K, value log is aboutg = 7 15%.4 0 higher.4, and than He hampering the determination of atmospheric parameters. 400 ± ± with increasing Teff. The O VI λλ1032/1038 resonance doublet =the0. spectroscopic25 0.1. The distance− error in derived the individual by Liebert metal et al. abundances(2013), which is ± . posesTo a mitigate problem. this At problem, the adopted we used abundance, SPY spectra, the innermost which have line a estimatedessentially not reflects to exceed the ratio0 5dex. of the effective temperatures of the cores are too deep and the line wings are slightly too shal- In Fig.3, we compare± the observed optical and near-UV H higher resolution ( 0.1Å), allowing us to trace the photospheric atmospheric model they used (93 235 K) and that of our model low. Interestingly, however, two subordinate O VI lines (4d–5 f and He lines with our final model. Although= nebular emission line profiles closer≈ toward the line core, because the nebular (105 000 K). It agrees with the value of d 870 250 pc derived and 4 f –5g transitions at 1123 and 1125 Å) are reproduced well linesby Frew mask et the al.(2016 Balmer) with line a centers, newly developed we note a statistical Balmer± line distance prob- emission lines are rather narrow. However, the signal-to-noise lem (Napiwotzki 1993), that is, the depth of innermost parts of ratioby our (S/ finalN) of model. these spectra They areis relatively too weak poor. in the In total,Teff = three 95 000 SPY K indicator for PNe. The distance to the WD given in the Gaia Data model (log g = 7.0) and too strong in the T = 115 000 K model theRelease observed 2 (1.36 photospheric+0.40 kpc) is profilesa factor of of two Hα largerand H thanβ is our not result, quite spectra (each consisting of three segmentseff covering the ranges 0.25 (log g = 7.4). matchedbut in agreement by the− models. within error Gianninas limits. et We al. conjecture (2010) were that unable the Gaia to 3400–4520Å, 4620–5600Å, and 5680–6640Å) were used and A weak F VI λ1139.50 line can be detected, allowing for a detectparallax this measurement problem because may beof the compromised lower resolution by the of strong their spec- emis- rectified. We abandoned co-addition because of the very differ- fluorine abundance measurement. The silicon abundance was tra.sion In lines order from to fit the these nearby lines, CEK, a lower which temperature could be would variable be re- in ent S/N of the single spectra in different wavelength regions. In derived from a fit to the Si IV λλ1122/1128 doublet. The blue strength, as was discussed in detail by Bond et al.(2016). the SPY spectra, too, no metal lines are detectable. 2 http://astro.uni-tuebingen.de/~TMAP A73, page 5 of9 Article number, page 4 of 9 A&A 616, A73 (2018) K. Werner et al.: The hot white dwarf in the peculiar binary nucleus of the planetary nebula EGB 6

EGB 6 FeVIII ArVII Si IV Fe VII 1.0 EGB 6 O VI relative flux relative flux 0.5 O IV O V O IV O IV O IV Ni VI 1056 1058 1060 1062 1064 1066 1068 1070 1072 1074 1076 1078 1080 1056 1058 1060 1062 1064 1066 1068 1070o 1072 1074 1076 1078 1080 o wavelength / A

Fig. 4. Detail of the FUSE spectrum and model for the white dwarf in EGEGBB6 6 from Fig.1.Fig.1.

areTable not 1. observed.Parameters Thefor the variations white dwarf in in log EGBg , 6.Teff , and He abun- HNFPHe CO SiSAr FeNi are not observed. The variations in log g , Teff , and He abun- 0 HNFPHe CO SiSAr FeNi iv 0 dance described in Sect.3.1 do not improve the C iv line fits. The same problem was encountered in our analysis of a hot The sameTeff problem/K was105 encountered 000 5000 in our analysis of a hot DO WD with1 similar2 temperature± and gravity (PG1034+001, DOlog WD(g withcm− similars− ) temperature7.4 0.4 and gravity (PG1034+001, 115000/7.0; Werner et al. 2017,± their Fig.1). The star is too hot 115000E/7.0;(B WernerV) et 0al..025 2017,0 their.005 Fig.1). The star is too hot to show the− C iii multiplet at 1175Å,± which is often seen in other to show the C iii 2multiplet at 1175Å, which is often seen in other -5 DAOlog WDs(nH andcm in) Tol26 (see20.3 Sect.3.4).0.4 -5 DAO WDs and in Tol26 (see Sect.3.4).± NitrogenM/M displays the0 N.58iv+multiplet0.12 at 922–924Å and the log (mass fraction) Nitrogen displays the N iv multiplet0.04 at 922–924Å and the log (mass fraction) − singlet atR 955/R Å. We also0. see026+ two0.021 Nv doublets at 1048Å and singlet at 955 Å. We also see two0.009 N doublets at 1048Å and 1050Å, which are partially blended+−300 by interstellar lines. The rel- 1050Å, whichL/L are partially blended72 48 by interstellar lines. The rel- 5 10 15 20 25 ative strength of the N iv and− N v lines is well matched by our 5 10 15 20 25 ative strength of the N iv and+ N1100v lines is well matched by our atomic number final model.d/pc Values of Teff 660as high400 as 115000K (at log g = 7.4) atomic number final model. Values of Teff as high− as 115000K (at log g = 7.4) or as low as 95000K1 (at log g+1=.267.0) can be strictly excluded or aslog low(tevol as 95000Kyr− ) (at3 log.60g0=.097.0) can be strictly excluded Fig. 5. Element abundances measured in the white dwarf in EGB 6 (red because the N iv lines are much− too weak or too strong, respec- Fig. 5. Element abundances measured in the white dwarf in EGB 6 (red becauselog(M the˙ /M N ivyrlines1) are much11.0+ too1.1 weak or too strong, respec- squares; see also Table 1). The blue line indicates solar abundances. − v 0.8 squares; see also Table1). The blue line indicates solar abundances. tively, while the N v line strengths− − are almost unaffected by the Teff variations. Teff variations.Abundances Ni/NH Xi [Xi] Prominent oxygen lines from three ionization stages are de- quired;3.4. The however, hot central this star would of Tolgive 26 too-deep Hγ and Hδ line cores ProminentH oxygen lines from 1 three ionization 0.74 stages 0.00 are de- tected. The O iv/O v line strength ratio is very sensitive against and, as mentioned above, would be in conflict with the metal line tected. TheHe O iv/O v line strength 0.085 ratio is very 0.25 sensitive0.00 against In Fig.6 we show the HST/COS spectrumii of Tol 26. Photo- Teff variations. While the O v lines hardly change within a range strengths. The central emission in Heii λ4686 is of nebular ori- Teff variations. While the O. v lines hardly4 change. within3 a range strengths. The central emission in He λ4686 is of nebular ori- C 3 4iv 10− 3 0 10− 0.10 sphericgin because lines its from height He, is C, not N, reachedO, and Fe by can the be emission identified. reversal of of 95000–110000K, the O iv× lines5 quickly× disappear4 with in- gin because its height is not reached by the emission reversal of N vi 1.9 10− 2.0 10− 0.54 theHe photosphericII λ1640 is line dominated profile. by strong nebular emission, but creasing Teff . The Ovi λλ1132× /11384 resonance× doublet3 − poses a the photospheric line profile. creasing TeOff . The O 41132.3 101138− resonance5.0 10 doublet− poses0.06 a we noteComparing the outer our wings results of with the photospheric the most recent absorption ones by profile. Gian- problem. At the adopted abundance,× 7 the innermost× line6 cores− are Comparing our results with the most recent ones by Gian- too deep andF the line wings1.1 are10− slightly1 too.5 shallow.10− Interes0.47 t- Shortwardninas et al. of (2010, 1180 Å,see the Sect.1), quality the of main the spectrum progress iswe somewhat achieved too deep and the line wings× are slightly5 too× shallow.4 Interest- ninas et al. (2010, see Sect.1), the main progress we achieved ingly, however,Si two subordinate3.9 10 O−vi lines8 (4d–5f.0 10 and− 4f–5g0.08 tran- pooris the but rather the tight C IV constraintdoubletat of the1169 eff Åective and thetemperature C III multiplet ( 5 % ingly, however, two subordinate× O vi7 lines (4d–5f× and6 4f–5g tran- is the rather tight constraint of the effective temperature (±5 % P 2.2 10− 5.0 10− 0.07 atcompared 1175 Å areto detectable.22%) by using The ionization C IV doublet balances might of be metals. filled± Inin sitions at 1123Å and 1125Å) are reproduced well by our final compared to ±22%) by using ionization balances of metals. In sitions at 1123Å and 1125Å). × are reproduced6 . × well4 by− our. final ± model. TheyS are too weak4 in3 the10T−eff = 95000K1 0 10 model− (log0 49g = addition,partiallyby the an error He inI λ the1168.67 surface airglow gravity line could (second be reduced order ffromrom model. They are too weak in× the Te6ff = 95000K× model5 − (log g = . . / 7.0) and tooAr strong in the 2T.e5ff =10115000K− 7 model.3 10 (log− g =0.007.4). λabout584.33).0.8 The dex to photospheric0.4dex. The C IV sameresonance holds for doublet the He/ isH abun- com- 7.0) and too strong in the Teff ×= 115000K model× (log g = 7.4). about ±0 8 dex to ±0 4dex. The same holds for the He H abun- Fe 8.5 10 5 3.5 10 3 0.46 dancepletely± ratio. filled As in by a consequence,± nebular emission. stellar We parameters note a photospheric can be det Ner-V A weak Fvi λ1139.50 line× can− be detected,× allowing− for a A weakNi F 1139.502. line3 10 can6 be detected,1.0 10 allowing4 0.17 for a minedresonance with doublet, higher precision but a potential (Sect.3.3). photospheric N IV λ1718 line fluorine abundancemeasurement.× The− silicon abundancewas× − de- is masked by a tenuous emission feature. We see a O IV multiplet rived from a fit to the Si iv λλ1122/1128 doublet. The blue com- λ ponentNotes. Mass, is blended radius, by luminosity, one of the post-AGB mentioned age relativelytevol, and mass-loss broad ph rateo- at3.2. 1340 Metal Å and abundances O V 1371. Their relative strengths can be used to ponent˙ is blended by one of the mentioned relatively broad pho- 3.2. Metal abundances M from evolutionaryvi tracks by Miller Bertolamiiv (2016). Abundances constrain the effective temperature. We identify a number of Fe V tosphericare given in O vi numberlines. ratios Another relative strong to hydrogen Si iv doublet (column at 2), 1067Å in mass is Welines, detected mostly lines in the from 1374–1430 ten metals Å region. in the FUSE spectrum (Figs. 1 blendedfractions by (column a strong 3), interstellar and logarithmic line. mass The phosphorusfractions relative abund toance solar v andTo 4). fit These the spectrum, lines are well we startedknown fromto use previous models analyses.with the com- Line wasvalue derived (column from 4; solar the abundancesP v λλ1118 from/1128 Asplund resonance et al. doublet. 2009). Error The vi lists,position from and which surface detailed gravity laboratory we adopted wavelengths for EGB6.and We tran adjustedsitions sulfurlimits for abundance the abundances was found are from0.5 dex the for S thevi λλ metals933/ and945 resonance0.4 dex for ± vi ± canTeff bebyusingobtained, the relative were presented O IV/O V byline Werner strengths. et al. The (2015, best theirfit is doublet,helium. as from two weaker subordinate S vi lines at 1000Å Tablesobtained 2–5) at T andeff = Werner85 000 et5000 al. (2017,K. At their that temperature, Tables 2 and the 3). C TheIII and 1118Å. The argon abundance was measured from the rather ± vii derivedλ1175 multiplet abundances (which are listed is not in present Table 1 in and EGB displayed 6 because in Fig of.5. its prominentand surface Ar gravityvii λ1063.55 could in line, principle whichis be blended constrained by the with wing He ofII iv an adjacent interstellar line. higherFromTeff carbon,) also fits two well. C iv Indoublets addition, are the seen Fe inV line the FUSE features spec- are anλ1640, adjacent but both interstellar quantities line. cannot be derived independently given trum,reproduced, namely so the that 3p–4d the C and and 3d–4f Fe abundances transitions are at the 1108Å same and as thatThe we strongestonly see the iron outermost lines stem part from of Fethevii line, followed wings. To by check Fe vi in EGB 6. However, we had to reduce the nitrogen and oxygen to what extent these uncertainties affect the determination of the 1169Å, respectively. In our final model with C = 0.003, the line and Fe viii lines (the only unblended Fe viii line is at 1148Å). abundances. We found N = 5 10 6 and O = 1.5 10 3, i.e., other atmospheric parameters (T , abundances of C, N, O, Fe), cores are not deep enough to match− the observation. A detaile− d The relative strengths of lineseff from the different ionization 40 and 3 times less than in EGB× 6. Error limits for× abundances we computed models varying log g ( 0.4 dex) and the He abun- inspection shows that the computed line profiles show central stages are sensitive temperature indicators. For example, at Teff = are estimated≈ to 0.5 dex. All other metal abundances cannot be dance (reduction from solar to 1/8 solar)± and we found thateff our emission cores. This is unrealistic, but despite some numeri- 95000K (and log g = 7.0) the Fe vii lines become too weak, assessed with the HST spectrum. The Si IV resonance doublet is results do not change within error limits, in particular, the N and cal testing concerning the temperature structure of the model, while the Fe vi lines become too strong. At Teff = 115000K (and obviously perturbed by interstellar absorption because the line O deficiencies remain. eff the reason remains unknown. Increasing the carbon abundance log g = 7.4) the Fe vi lines are too weak, while the Fe viii lines depths cannot be matched by a photospheric model even with Given the uncertainty in the He abundance, the ques- mitigates the problem in the line cores, but leads to too strong are too strong. The nickel abundance is derived from a multitude an unrealistically high Si abundance. The helium abundance tion arises whether Tol 26 is a DAOZ at all. Could it be a line wings and to the appearance of optical C iv lines, which of relatively weak Ni vi lines. A73, page 6 of9 Article number, page 5 of 9 K. Werner et al.: The hot white dwarf in the peculiar binary nucleus of the planetary nebula EGB 6 A&A proofs: manuscript no. aa

Tol 26

N V 1

O IV C III C IV Lyα

0 1150 1160 1170 1180 1190 1200 1210 1220 1230 1240 1250 1260 1270 1280 1290

Fe VI Si IV

1.0

O IV Fe V Fe V 0.6 O V relative flux 1300 1310 1320 1330 1340 1350 1360 1370 1380 1390 1400 1410 1420 1430

1.0

C IV 0.5 1480 1490 1500 1510 1520 1530 1540 1550 1560 1570 1580 1590

1.0

0.6 HeII N IV

1600 1610 1620 1630 1640 1650 1660 1670 1680 1690 1700 1710 1720 1730 1740 1750 1760 o wavelength / A

Fig. 6. HST/COSHST/COS spectrum of the white dwarf in Tol Tol 26 (gray) (gray) compared compared to to a a p photospherichotospheric model model spectrum spectrum (red; (red;TTeeffff==85 85 000 000 K, K, loglog gg==7.4) 7.4) with abundances as given in Sect. 3.4 3.4.. The blue graph is the phphotosphericotospheric model attenuated by interstellar lines.Promin Prominentent photospheric lines are identified.identified. Absorption Absorption profiles profiles of of the the C CIViv resonanceresonance doublet doublet and and He HeiiIIλλ16401640 Å Å (except (except for for the far wings) are masked by nebular emissiemissionon lines.

He-dominatedOther iron white group dwarf elements (i.e., spectral were not type detected, DOAZ) and or even we lutionaryuncertain tracks (Teff = 105 by Miller000 5000 BertolamiK, log (2016,g = 7.4 Fig.7).0.4). We A mass find +0+.012.12 ± +300 ± completelycan thus exclude devoid extremeof hydrogen overabundances (spectral type that DOZ)? were Probably seen in Mof /MM/M==00.58.58 . andwasL derived./ L = 72 The abundances. The errors of are H, dom- He, ⊙ 0 004.04 ⊙ 48 not.other When DAOs, we namelyremove hydrogenHS2115+ from1148 our (Werner model, et the al. He 2018)II λ1640 and inatedand ten by more the species uncertainty−− are solar. in log g . The− post-AGB age is about wingsBD 22 become◦3467, the so exciting strong and star broadof Abell that 35 it (Ziegler is not et compatible al. 2012). 4000yr;We conclude however, that the element mass uncertainty diffusion by causes gravitational a large possib settlingle − +1.26 withLikewise, the observation, we found no unless hint of the trans-iron surface gravity group iselements reduced that beloware error:and radiative log(tevol/ levitationyr) = 3.60 is0.09 not. The acting mass-loss in the rates photosphere. along the evo- The logpresentg = 7 in in strongly order to oversolar diminish the amounts strength in ofseveral the line hot wings. DO WDs lutionprobable assumed reason by is Miller that the− Bertolami relatively (2016) luminous imply star log( ( M100˙ /ML /)yr) is ≈ ⊙ (Hoyer et al. 2018, and references therein). =losing11.0 mass+1.1. by a radiation-driven wind at a rate on the order of −11 0.8 1 To summarize, the abundances of hydrogen, helium, and ten 10−TheM− spectroscopicyr− . Compared distance to otherd was DAOZ found WDs, bycomparing this is remark- the 4. Summary and discussion more identified species in the white dwarf are, within error lim- dereddenedable because visual they magnitude usually showV0 with strong the deviations respective from model solar at- 4.1.its, compatible EGB 6 with solar values. mosphereabundance flux, patterns resulting due in to the diffusion relation (e.g., Good et al. 2005; Ziegler et al. 2012; Werner et al. 2018). The sample of DAOZs We performed a spectral analysis of the hot WD in the and related hot progenitors4 analyzed by 0Ziegler.4V0 log g(2012) con- 3.3. Stellar parameters, distance d[pc] = 7.11 10 Hν M 10 − , PN EGB 6. The effective temperature could be tightly tains objects with temperature× and· gravity· similar to EGB 6. In 3 2 1 1 constrained,Stellar parameters while are the estimated surface from gravity the remains comparison relatively of the whereparticular,Hν the= central1.60 star10− oferg Abellp cm 31− s (−TeffHz=−114is 000 the Eddington10 000 K, × ± determined atmospheric parameters Teff and log g with evo- flux of the model at 5400Å, and M is the stellar mass in M . ⊙ A73, page 7 of9 Article number, page 6 of 9 K. Werner et al.: The hot white dwarf in the peculiar binary nucleus of the planetary nebula EGB 6 A&A 616, A73 (2018)

Table 1. Parameters for the white dwarf in EGB 6.a 10 yr; see, e.g., the compilation maintained by D. Jones3) than 0.532 that≈ of EGB 6 (larger than 500 yr); however a number of wide, 0.566 ≈ Teff / K 105000 5000 7 visual binaries are known as well (separations of 100 to a few 2 ± ≈ log(g/cm/s ) 7.4 0.4 0.583 thousand AU; Ciardullo et al. 1999). ± ) Liebert et al.(2013) considered an alternative explanation E(B V) 0.025 0.005 2 − ± 0.616 for the origin of the compact emission knot. They put forward log(n /cm 2) 20.3 0.4 H − ± the idea that it is the region where the winds from the WD +0.12 and the supposed dM companion collide. Their “duelling-winds” M / M 0.58 0.04 ⊙ − model was considered unlikely by Bond et al.(2016) since they R / R 0.026+0.021 EGB 6 ⊙ 0.009 log (g / cm/s 0.706 discovered that the companion is not an exposed dM star but +− L / L 72 300 a dust-enshrouded low-luminosity star. In view of our result ⊙ 48 −+1100 that the WD is rather young and luminous, it might be worth d / pc 660 400 − 8 reconsidering the colliding-wind hypothesis. From an ioniza- log(t /yr) 3.60+1.26 evol 0.09 tion model for the CEK, Liebert et al.(2013) estimated that ˙ −+1.1 log(M/M /yr) 11.0 0.8 5.2 5.0 4.8 a WD wind with a rather high mass-loss rate on the order ⊙ − − 9 1 log (T /K) of 10− M yr− would be required, but they questioned how a eff abundances Ni/NH Xi [Xi] hydrogen-rich WD could maintain such a high rate. g H 1 0.74 0.00 Fig. 7. PositionPosition of of the the white white dwarf dwarf in inEGB EGB6 6 with with error error bars bars in the ing the–Teff– We have shown that the WD has a solar metallicity; there- He 0.085 0.25 0.00 Tdiagrameff diagram (blue (blue cross). cross). Red, Red, green, green, and and black black symbols symbols are are DAO DAO WWDsDs fore, a radiation-driven wind can be sustained with a mass-loss 4 3 from the analyses of Gianninas et al. al.( (2010),2010), Ziegler Ziegler(2012 (2012)), and, and Werner Werner rate of up to 10 10 M yr 1, and thus almost reaches the required C 3.4 10− 3.0 10− 0.10 et al. (2018), respectively. Black lines are evolutionary tracks of post- − − × 5 × 4 et al.(2018), respectively. Black lines are evolutionary tracks of post- value. However, we also point out that there is no spectroscopic N 1.9 10− 2.0 10− 0.54 AGB remnants with masses as labeled (in M , from Miller Bertolami × 4 × 3 − ⊙ hint of mass-loss; in particular, we would expect the O VI res- O 4.3 10− 5.0 10− 0.06 2016,2016; metallicity Z = 0..0101).). × 7 × 6 − onance doublet to be most susceptible to wind effects, but the F 1.1 10− 1.5 10− 0.47 × 5 × 4 line centers are not shifted relative to other photospheric lines Si 3.9 10− 8.0 10− 0.08 × 7 × 6 log g = 7.4 0.3) has a roughly solar abundance pattern, and the profiles are symmetric. Investigations of H-rich central P 2.2 10− 5.0 10− 0.07 plet at 1175ű are detectable. The C iv doublet might be filled 9 1 × 6 × 4 − while slightly cooler objects like the central star of Abell 35 stars showed that the rates are usually larger than 10− M yr− . S 4.3 10− 1.0 10− 0.49 in partially by an He i λ1168.67g airglow. line. (second order from × 6 × 5 − (Teff = 80 000 10 000 K, log = 7 2 0 3) show extreme devi- The lowest rate derived so far, and thus constraining the thresh- Ar 2.5 10− 7.3 10− 0.00 λations584.33).The from it± photospheric (Ziegler et Cal.iv 2012resonance).± Obviously, doublet isstellar completely winds × 5 × 3 old for the spectroscopic proof of existing winds, is that of Fe 8.5 10− 3.5 10− 0.46 filled in by nebular emission. We note a photospheric N v reso- 10 1 × × in cooling DAO WDs effectively prevent element diffusion as NGC 1360 with 10− M yr− (Herald & Bianchi 2011) using . 6 . 4 nance doublet, but a potential photospheric N iv λ1718 line is ≈ Ni 2 3 10− 1 0 10− 0.17 long as Teff is higher than 100 000 K. Our result for Tol 26 very weak deviations of the O VI resonance doublet line pro- × × masked by a tenuous emission≈ feature. We see a O iv multiplet (Teff 85 000 K; O and N depletion) confirms this conclusion. files from hydrostatic-atmosphere profiles. Detailed expanding (a) ≈ v Notes. Mass, radius, luminosity, post-AGB age tevol, and mass-loss atSimilarly, 1340Å the and correlation O λ1371. between Their relative helium strengths abundance can and be effec- used model-atmosphere computations tailored to EGB 6 are required rate M˙ from evolutionary tracks by Miller Bertolami (2016). Abun- totive constrain temperature the e foundffective by temperature.Gianninas etWe al.(2010 identify) in their a number sample of to determine the mass-loss rate that would be necessary to leave dances are given in number ratios relative to hydrogen (column 2), in Feofv DAOslines, was mostly interpreted in the 1374–1430Å as an indication region. of how the weaken- a detectable spectroscopic signature. mass fractions (column 3), and logarithmic mass fractions relative to ingTo stellar fit the wind spectrum, along the we WD started cooling to use sequence modelswith can no the longer com- solar value (column 4; solar abundances from Asplund et al. 2009). Er- maintain helium in the atmosphere of a white dwarf. ror limits for the abundances are 0.5 dex for the metals and 0.4 dex position and surface gravity we adopted for EGB6. We adjusted 4.2. Tol 26 ± ± The theoretical mass-loss rate from the hot WD, for helium. Teff by using the relative O iv/O v line strengths. The best fit ˙ / 1 . +1.1 islog obtained(M M yr at− T)eff == 8500011 0 0.8,5000K. is orders At of that magnitudes temperature, larger the We analyzed a UV spectrum of Tol 26. We found, as for EGB 6, − − Cthaniii λ required1175 multiplet to inhibit (which Bondi–Hoyle± is not present accretion in EGB6 of becauseinterstellar of that the hot nucleus of Tol 26 is a DAOZ WD, but cooler material (MacDonald 1992). We therefore do notv expect that (Teff 85 000 K). Helium abundance and surface gravity are its higher Teff ) also fits well. In addition, the Fe line features ≈ For our WD we have V = 15.998 (Bond et al. 2016). From areany reproduced, circumstellar so materialthat the C from and Fe the abundances PN is accreted are the unless same uncertain with the observations currently at hand. We measured E(B V) = 0.025 we derived the visual extinction using the asthe in accretion EGB6. However, flow is reorganized we had to reduce into a the disk, nitrogen due to and angular oxy- the abundances of C, N, O, and Fe. Significant depletion of nitro- − . . . 6 3 standard relation AV = 3 1E(B V) = 0 077, hence V0 = 15 92. genmomentum. abundances. We found N = 5 10− and O = 1.5 10− , gen and oxygen was revealed, which can be interpreted as an +1100 − × × We found d = 660 400 pc. This value is about 15% higher than i.e.,Concerning 40 and 3 times the outer less nebula,than in EGB6.Liebert Error et al.( limits1989) for pointed abun- indication that the gravitational settling of elements affects the the spectroscopic distance− derived by Liebert et al. (2013),which dancesout a discrepancy are estimated≈ between to 0.5dex. their All estimations other metal of abundances the post-AGB can- abundance pattern in the photosphere, in contrast to the WD in 5 essentially reflects the ratio of the effective temperatures of the notage be of assessed the WD with (several the HST10 yr; spectrum. based on The their Si iv spectralresonance analysis dou- EGB 6. Our results confirm that, with regard to the hot nucleus atmospheric model they used (93235K) and that of our model bletand evolutionaryis obviously perturbed tracks by Koester by interstellar & Schönberner absorption 1986 bec)ause and the and the compact nebula emission region, Tol 26 is similar to (105000K). It agrees with the value of d = 870 250pc derived kinematical age of the nebula (20–50 kyr, based on their derived EGB 6. It remains to be investigated whether the central star of ± line depths cannot be matched by a photospheric model even by Frew et al. (2016) with a newly developed statistical distance withspectroscopic an unrealistically distance of high 460 Si pc abundance. and assuming The a PN helium expansion abun- Tol 26 also has a low-luminosity companion. 1 indicatorfor PNe. The distance to the WD given in the Gaia Data dancerate of and 20–50 surface km gravity s− ). Our could results in principle indicate be a constrained much youngerwith +0.40 1 . +1.26 Acknowledgements. We thank R. Napiwotzki for putting the SPY survey spec- Release 2 (1.36 0.25 kpc) is a factor of two larger than our re- HeWDii λ age,1640,log but(tevol bothyr− quantities= 3 60 cannot0.09), mainly be derived due independently to the higher sult, but in agreement− within error limits. We conjecture that the effective temperature (105 000− K compared to 70 000 K). Tak- tra at our disposal. The TMAD tool (http://astro.uni-tuebingen.de/ given that we only see the outermost part of the line wings. To ~TMAD) used for this paper was constructed as part of the activities of the Gaia parallax measurement may be compromised by the strong checking the to PN what angular extent extension these uncertainties (130), the a spectroscopicffect the determination distance 1 German Astrophysical Virtual Observatory. Some of the data presented in emission lines from the nearby CEK, which could be variable in ofderived the other by us atmospheric (660 pc), and parameters the expansion (Teff , abundances velocity (38 of km C, s− N,; this paper were obtained from the Mikulski Archive for Space Telescopes strength, as was discussed in detail by Bond et al. (2016). O,Hippelein Fe), we & computed Weinberger models 1990 varying), we obtain log g ( a0 kinematical.4dex) and age the (MAST). STScI is operated by the Association of Universities for Research Heof 15 abundance kyr, which (reduction is in agreement from solar with to 1/ the8 solar) post-AGB± and we age found of in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST the WD. for non-HST data was provided by the NASA Office of Space Science via 3.4. The hot central star of Tol26 that our results do not change within error limits, in particular, grant NNX09AF08G and by other grants and contracts. This research has made the NIf theand material O deficiencies around remain. the companion was captured at the PN use of NASA’s Astrophysics Data System and the SIMBAD database, oper- ejection 15 kyr ago, and if the disk lifetime is not much longer, In Fig.6 we show the HST/COS spectrum of Tol26. Photo- Given the uncertainty in the He abundance, the question ated at CDS, Strasbourg, France. Funding for SDSS-III has been provided by then this would explain why the EGB 6 phenomenon is uncom- the Alfred P. Sloan Foundation, the Participating Institutions, the National Sci- spheric lines from He, C, N, O, and Fe can be identified. arises whether Tol26 is a DAOZ at all. Could it be a He- mon. Also, the large distance of the companion from the primary ence Foundation, and the U.S. Department of Energy Office of Science. This He ii λ1640 is dominated by strong nebular emission, but dominated white dwarf (i.e., spectral type DOAZ) or even com- work has made use of data from the European Space Agency (ESA) mis- (at least 100 AU) prevents a fast evaporation of the disk mate- we note the outer wings of the photospheric absorption pro- pletely devoid of hydrogen (spectral type DOZ)? Probably not. sion Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia rial by UV≈ radiation from the primary. Measured orbital periods file. Shortward of 1180Å, the quality of the spectrum is some- When we remove hydrogen from our model, the He ii λ1640 of binary central stars are all much shorter (from hours up to 3 http://drdjones.net/bCSPN what poor but the C iv doublet at 1169Å and the C iii multi- wings become so strong and broad that it is not compatible A73, page 8 of9 Article number, page 7 of 9 K. Werner et al.: The hot white dwarf in the peculiar binary nucleus of the planetary nebula EGB 6

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