arXiv:0901.4553v1 [astro-ph.SR] 28 Jan 2009 n h ats tla id r bevdi WO in observed are abundance winds C+O stellar largest fastest The the and stages 1983). These evolutionary al. et to (Conti carbon, respectively. correspond nitrogen, types ac- lines, spectral WO, of emission and dominance oxygen WC, and the WN, to sorted are cording subclasses: spectra very three WR The and into winds. composition stellar chemical strong peculiar highly Introduction 1. icvr fXryEiso rmteWl-ae trW 4 o 142 WR Wolf-Rayet the from Emission X-ray of Discovery eateto hsc n srnm,Es ense tt U State Tennessee East Astronomy, and Physics of Department tr fteWl-ae W)tp aea have type (WR) Wolf-Rayet the of eateto srnm,Uiest fIlni,10 West 1002 Illinois, of University Astronomy, of Department R12wt the with 142 WR -as hc a edet togwn oiain idcupn,o n stars or clumping, wind headings: Subject ionization, rela be wind must strong 142 rotation. to WR rapid of due from wind be the can unusua that the which implies by X-rays, detection indicated binar Our as fast, a magnetic lines. extremely for emission to rotate hints refer to must no seems explanation 142 are WR speculative There ins albeit intrinsic remaining, radiation. the only observed mecha to the this that due of conclude shocks hardness we considerations embedded Co qualitative MK. to From 100 attributed driving. about is of temperature winds plasma stellar a suggests spectrum ray L asv tri eyavne vltoaysae hr before short stage, evolutionary advanced very γ a in star massive n eedptu -a ore within sources X-ray serendipitous any rybrt hsi h rtdtcino -a msinfo WO-t a from emission X-ray of detection first the is This burst. -ray X erpr h icvr fwa e adXryeiso rmteW the from emission X-ray hard yet weak of discovery the report We 7 = nttt o hsc n srnm,Uiest osa,1 Potsdam, University Astronomy, and Physics for Institute × 10 30 r s erg tr:wns uflw tr:Wl-ae tr:individu stars: — Wolf-Rayet stars: — outflows winds, stars: XMM-Newton − 1 .M sioa .R aan .Feldmeier A. Hamann, W.-R. Oskinova, M. L. hc osiue only constitutes which , [email protected] -a eecp.Bigo pcrlsbyeW2 R12i a is 142 WR WO2, subtype spectral of Being telescope. X-ray xgnsubtype oxygen ABSTRACT ≈ .H Chu Y.-H. .Ignace R. 1 ′′ ∼ < fW 4.W 4 a nXrylmnst of luminosity X-ray an has 142 WR 142. WR of 1 10 − 8 tg famsiesa ro oisepoinas explosion its to or prior Ic star evolutionary type massive final a a the of represent stage that stars, type ra ee aan(05.I hi simulation by their presented In was wind (2005). WR Hamann a Gr¨afener hydrody- & for first The model winds. namical for WR opacity of iron acceleration the of the importance (1996) the Schaerer out pointed winds. line-driven of derstanding 2005). fisblmti uioiy h adX- hard The luminosity. bolometric its of h Rsascniu ocalneorun- our challenge to continue stars WR The iest,JhsnCt,T 71,USA 37614, TN City, Johnson niversity, re tet raa L681 USA 61801, IL Urbana, Street, Green t xlso sasproaor supernova a as explosion its 46Ptdm Germany Potsdam, 4476 l on rfie fisoptical its of profiles round lly mny -a msinfrom emission X-ray mmonly, opno.Teeoethe Therefore companion. y imcno con o the for account cannot nism ciiy osbyrelated, Possibly activity. aiiyo h radiation the of tability p tr erl out rule We star. ype nshrclgeometry on-spherical l(R12 X-rays: — 142) (WR al l-ae W)star (WR) olf-Rayet ieytasaetto transparent tively γ rybrt(ish tal. et (Hirschi burst -ray f the mass loss is initiated at high optical depth by cations of the XMM-Newton detection of WR142 the so-called “iron bump” in the opacity. It was are discussed in Sect. 4. thus demonstrated that WR-type winds can be driven by radiation pressure. 2. The WO-type Star WR142 It has long been known that line-driven winds WR142, also named Sand5 and St3, has a are subject to an instability that can lead to spectrum characteristic for the spectral subtype strong shocks (Lucy & White 1980). These WO2 (Barlow & Hummer 1982; Kingsburgh et al. shocks are thought to explain the X-ray emis- 1994). The optical spectrum of WR142 was dis- sions from O star winds, as predicted by time- cussed by Polcaro et al. (1997), who also noticed dependent hydrodynamic modeling (Owocki et al. some line variability. 1988; Feldmeier et al. 1997) and largely confirmed by observations (Kramer et al. 2003; Oskinova et al. Figure 1 shows the spectral energy distribution 2006; Zhekov & Palla 2007; Waldron & Cassinelli (SED) of WR 142 together with model calculated 2007). The growth of instability in WR winds was with the Potsdam Wolf-Rayet (PoWR) model at- investigated by Gayley & Owocki (1995). They mosphere code (Gr¨afener et al. 2002). Photomet- found that despite of damping effects due to the ric IR measurements are plotted together with multi-line scattering, the instability remains effec- our optical spectrum and the Spitzer IRS mid- tive. Therefore, X-ray emission from wind shocks IR spectrum. We have adopted parameters typ- could, in principle, be expected in WR winds, a ical for a WO star: stellar temperature T∗ = conjecture that has not been yet tested by time- 150kK, “transformed radius” (cf. Gr¨afener et al. dependent hydrodynamic simulations. 2002) Rt = 2R⊙, and a composition of 40% carbon, 30% oxygen and 30% helium (by mass). Significant observational effort has been made WR142 is assumed to be a member of the open to study the X-ray emission of WR stars. White & Long cluster Berkeley 87 at a distance of d = 1.23kpc (1986), Pollock (1987), Pollock et al. (1995), (Turner et al. 2006). Based on this preliminary Oskinova (2005) presented X-ray observations model, the fit of the photometric observations of Galactic WR stars. A survey of X-ray emis- requires an interstellar reddening of EB−V = sion from WR stars in the 1.7mag and a stellar luminosity of log Lbol/L⊙ = was conducted by Guerrero & Chu (2008a,b). 5.35, implying a stellar radius of only R∗ = Ignace et al. (2000) and Oskinova (2005) demon- 0.6 R⊙. The corresponding mass-loss rate is about strated that X-ray properties of single WR stars −5.1 −1 10 M⊙yr for an adopted microclumping vol- differ from those of O stars. Whereas O stars dis- ume filling factor of 0.1. play a trend in which the ratio of the X-ray to the bolometric luminosity LX/Lbol has a typical value The adopted model does not provide an en- of 10−7 (Long & White 1980; Berghoefer et al. tirely satisfactory fit to the line spectra. For ex- vi 1997; Sana et al. 2006), this trend is not observed ample, the model does not match the huge O ˚ in the case of WR stars. emission at 3811, 3834 A, a problem also experi- enced by Crowther et al. (2000) while reproducing Observations with the XMM-Newton and Chan- these lines for Sand 2 with the cmfgen code. We dra X-ray telescopes established that some bona also fitted the SED shown in Fig. 1 with a model fide single WN stars are X-ray active (Skinner et al. that has a mass-loss rate lower by a factor of two, 2002a,b; Ignace et al. 2003; Oskinova 2005), while −5.4 −1 10 M⊙yr , and higher bolometric luminosity, others are apparently not (Oskinova 2005; Gosset et al. log Lbol/L⊙ = 5.65. This model fits the SED in 2005). Oskinova et al. (2003) found that no sin- Fig. 1 equally well. Figure 2 shows the radius of gle WC star had been conclusively detected at unity optical depth plotted as function of wave- X-ray energies, a result that continues to hold. length in the X-ray range for both models. The Among the WR subclasses, only the WO-type stronger wind is opaque even to hard X-rays, but stars have not been observed in X-rays so far. In the thinner wind is largely transparent, because this Letter we present the XMM-Newton observa- its higher ionization reduces the X-ray absorbing tions of the closest WO type star WR142. The ions. The same could happen in denser, but hotter star is introduced in Sect. 2, and its XMM-Newton models. Thus our ability to predict the influence observations are described in Sect. 3. The impli-

2 of wind photo-absorption is somewhat limited, ow- 12.0keV band are (1.89 ± 0.34) × 10−3 cs−1 for ing to ambiguities in the ionization state of metals EPIC MOS1+2 and (3.80 ± 0.84) × 10−3 cs−1 for and uncertainty in the mass-loss rate. EPIC PN cameras. Assuming a two-temperature The profiles of the emission lines in the spec- thermal plasma model (kT1 = 0.3 keV,kT2 = trum of WR 142 are very broad. Assuming that 10 keV), the observed X-ray flux of WR 142 is −14 −1 −2 the line widths correspond to the wind termi- FX = 4 ± 2 × 10 ergs cm . The redden- −1 nal velocity, the velocity of v∞ ≈ 5500kms ing towards WR 142 is known from the analysis would be deduced. However, the profile shapes of its optical spectrum, and the distance is known of almost all lines are much more round than the from its cluster membership. The X-ray luminos- 30 −1 roughly Gaussian shapes usually seen in WR spec- ity of WR 142 is thus LX ≈ 7 × 10 ergs , or tra. Such round profiles cannot be reproduced log LX/Lbol ≈−8. < ′′ by the standard models. It is tempting to repro- The angular resolution of XMM-Newton is ∼ 6 . duce these profiles by convolution with the semi- To exclude the potential confusion with a source ellipse for rotational broadening, albeit rotating in close vicinity of WR 142, we inspected opti- stellar winds certainly require a more sophisti- cal and infra-red images with higher angular res- cated treatment which has not been accomplished olution. According to the USNO-B1.0 catalog yet. If rotation is the cause for the round pro- (Monet et al. 2003), the closest object to WR142 files, the projected rotation speed must be com- is located 8′′ away. The optical monitor (OM) on parable to v∞, i.e. the star would rotate near board of XMM-Newton provides images with an < −1 ′′ to break-up with vrot sin i∼4000kms (see inset angular resolution of ≈ 1 . The OM images in in Fig. 1). Interestingly, a similar suggestion has four filters did not detect any objects closer than been made for the hottest and the most compact 8′′ around WR142. The Spitzer space telescope Galactic WN star, WR 2, (T∗ = 140kK, R∗ = has observed the Berkeley87 cluster in the infra- −1 < ′′ 0.89 R⊙, vrot sin i ≈ 2000km s ) (Hamann et al. red (IR). With an angular resolution of ∼ 1 , the 2006). Spitzer IRAC camera took images in four chan- nels. No point sources within 8′′ from WR142 3. Observations were detected in any of the IRAC channels. It is extremely unlikely that any X-ray source is not de- WR142 was observed by XMM-Newton during tectable in either the optical or IR. We therefore two consecutive satellite orbits (ObsId 0550220101 rule out any serendipitous X-ray sources within and ObsId0550220201). The data were merged ≈ 1′′ of WR142 that might get confused as our and analyzed using the latest versions of software target. sas 8.0.0. After the high background level time In addition, there is no evidence to suggest intervals have been rejected, the combined expo- that the observed X-rays originate from a wind sure time of all detectors was ≈ 100ks; an EPIC blown bubble around WR 142. Although theoreti- image of WR 142 and its surroundings is shown cally expected, there is a dearth of WR stars with in Fig. 3. WR 142 is detected with a 5σ confi- detected diffuse X-ray emission from their wind- dence level in all XMM-Newton EPIC detectors blown bubbles (Chu et al. 2003; Wrigge et al. using standard source-detection algorithms. We 2005). The only two detected hot bubbles show define a “hardness ratio HR” as HR = (Nhard − −1 a limb-brightened morphology and are extended Nsoft)(Nhard + Nsoft) , where Nsoft is the number on the scale of . On the contrary, we con- of counts in the 0.25-2.0 keV band and Nhard in the fidently detect a point source at the position of 2.0-12.0 keV band. For WR 142 we find HR≈ 0.57. WR 142. For comparison, the hardness ratio of the WN star The spectra of WR142 were extracted from the WR 1, which has a reddening similar to WR 142, is ′′ HR(WR 1)≈ −0.9, while more reddened WR110 15 region. The small number of counts obtained has HR(WR 110)≈ −0.4 and less reddened WR6 in our observation does not allow for quantitative has HR(WR 6)≈−0.6. Thus X-ray emission from spectral analysis, but interesting conclusions can WR 142 is the hardest among putatively single be made about the gross energy distribution. Fig- WR stars. The count rates in the 0.25keV- ure4 shows the X-ray spectrum of WR142 before a background subtraction, together with the spec-

3 1 trum of a nearby background region normalized to µ ≈ 2 for ionized H, while for fully ionized car- 2 the same area. The emission from WR142 domi- bon, µ ≈ 3 . Thus, differences in wind molecular nates over the background at 2−7 A˚ (1.7−6 keV), weight between O and WR stars cannot strongly while longwards it dives below the background un- affect the temperatures of the shocked plasma in til about ≈ 20 A˚ (0.6 keV) where it rises above the stars of these types. background again. The presence of “a dip” be- Note that Eq. (1) gives an upper limit to the tween ≈ 7 − 12 A˚ (0.8 − 1 keV) is also seen in the temperature at the shock. Radiative cooling will background-subtracted PN and MOS-2 spectra restrict the hottest temperatures to a limited part shown in Fig.5. The appearance of the dip may be of the total emission measure. Moreover, some due to unresolved strong emission lines. However, fraction of the energy will be consumed by ion- the wavelength of the dip coincides well with the ization processes. To roughly estimate the energy K-shell edges of oxygen (Verner & Yakovlev 1995), required to ionize WO wind material in the post- where our models predict an absorp- shock zone we consider oxygen and assume that tion maximum (see Fig. 2). Therefore it is tempt- the leading ion in the pre-shock material is O vi, ing to attribute the observed dip to the oxygen while in the post-shock plasma it is O ix. AWO K-shell absorption – to confirm this identification, wind may contain 30% of oxygen. The ioniza- data of higher quality will be required. Assum- tion potentials are IP(O vi)≈ 0.14keV, IP(O vii)≈ ing the presence of K-shell of oxygen, the thermal 0.74keV, and IP(O viii)≈ 0.87 keV (Cox 2000). plasma model fit to our low S/N data indicates The specific energy for ionization is ǫ = ΣNiIPi, −1 presence of plasma with temperatures spanning where Ni = Xi(AimH) with Xi being the mass from 1MK to 100MK. fraction of the element, Ai its atomic weight, and mH the atomic mass unit. Inserting these num- 4. On The Origin of X-ray Emission from bers, the full ionization of oxygen requires ǫ ≈ WR 142 3 × 1013 ergg−1. This is small compared to the ki- netic energy of the wind (5 × 1015 ergg−1 for typi- The confident detection of weak hard X-rays 1 cal 1000kms− ). Thus the ionization is not a sig- from WR 142 prompts us to re-consider the pre- nificant cooling process even in metal-rich winds. vious non-detections of X-ray emission from WR In the time-dependent hydrodynamic simula- stars. The X-ray luminosity of WR 142 of LX ≈ 7×1030 ergs−1 is comparable to present upper lim- tions by Feldmeier et al. (1997) the velocity jump its for the non-detections (Oskinova et al. 2003; U depends on the ratio between the period of Oskinova 2005; Gosset et al. 2005; Skinner et al. the perturbations at the wind base, Tc, and the flow time, Tflow = R∗/v∞. The former was esti- 2006). Perhaps some WR stars that are yet un- −1 detected in X-rays may be weak sources similar mated by the acoustic cut-off period, Tc = aH , to WR 142. This raises the question, what is the where a is thermal speed and H is the pressure mechanism responsible for the generation of X- scale height. Assuming that the perturbation are rays in WR stars? Usual suspects are wind shocks seeded in the hydrostatic layers within the star, ∝ −1 and magnetic fields. Tc/Tflow a vesc. The thermal speed is lower in hydrogen-deficient WR atmospheres, while vesc The strong-shock condition predicts the peak can be larger than in O stars, as it is the case temperature in the shocked gas, TX, as for our program star WR142. Therefore, the ve-

max 3 2 locity jumps in WR wind shocks are expected to kTX = µmHU , (1) 16 be rather smaller than larger, compared to O-star winds. where U is the velocity with which the gas rams These qualitative considerations indicate that into the shock (i.e., pre-shock velocity relative to if the same shock mechanism were in operation in the shock), and µ is the mean molecular weight per WR and O star winds, the X-ray spectra from WR post-shock particle in the gas. As the shock tem- stars would be softer than those of O stars. How- peratures are very high, abundant species will be ever, observations show the contrary – the handful nearly entirely ionized. Therefore, µ depends only of single WR stars with available spectral mea- weakly on chemical compositions. For example, surements all show X-ray emission harder than

4 typically found in O stars (Skinner et al. 2002a,b; the mass-loss (a poorly constrained parameter) Ignace et al. 2003). This trend is confirmed by our by a factor of only two and/or a higher effec- new data on WR142. tive stellar temperature result in a higher degree Relatively hard X-ray emission can occur in a of wind ionization. In this case a fraction of X- binary system; however, optical spectra of WR142 rays formed deep in the wind could freely escape. show no contamination from an early-type com- Wind clumping can further reduce wind attenua- panion. A low-mass coronal companion to this ini- tion (Feldmeier et al. 2003). tially very massive star is extremely unlikely (Lucy There is suggestive evidence that WR 142 is 2006). a fast rotator, as is the case for the WN2 star Babel & Montmerle (1997) proposed that if a WR 2 (see Sect. 2). The detection of strong and 32 −1 stellar wind is magnetically confined it can be hard (log LX ≈ 10 ergs , TX ≫ 10MK) X-rays strongly heated. The magnetic field locally dom- in WR 2 (Skinner et al. 2008) likely indicates that inates the bulk motions if the magnetic energy a similar mechanism operates in these hot, com- density exceeds the wind kinetic energy density, pact, fast rotating stars. The lower X-ray lu- 2 2 B /µ0 > ρv . At characteristic distance of 1 R∗ minosity of WR142 compared to WR2 can re- from the photosphere, the velocity is v ≈ 0.5v∞ ≈ flect the higher absorption of X-rays in its denser 2000kms−1, density ρ = M/˙ 4πv(r)2r ≈ 1 × wind. Fast rotation can lead to a slow equatorial −10 −3 10 g cm . Therefore, B(r = 2R∗) > 7 kG wind with enhanced density, possibly forming an is required to control the WR142 wind. For a out-flowing disk, and a faster, thinner polar wind dipole field, this implies that the magnetic field (Ignace et al. 1996). If the round shape of the line strength at the surface is larger than 50kG. A profiles found in both, WR 142 and WR 2, is due to field of such strength is not unrealistic, since the rotational broadening, the inclination angle must radius of WR 142 is ∼ 50 times smaller than that be large. of OB supergiants, where surface fields of 100G Summarizing, we have reported the detection of have been observed (Bouret et al. 2008). One may weak, but hard X-rays from WR142 and speculate speculate that even less strong magnetic fields may that their origin is connected with magnetic activ- lead to hard X-rays due to some type of reconnec- ity of this far-evolved, compact WR star. The spe- tion and heating processes as invoked to explain cific shape of its optical line profiles signals that the high X-ray temperatures observed in O type WR142 may be rotating at nearly break-up ve- stars (Waldron & Cassinelli 2007). locity. In this case the spherical symmetry will It is difficult to test where in the wind the ob- be broken, potentially affecting the X-ray produc- served X-rays have been produced. Since no high- tion and absorption. Observations of better qual- resolution X-ray spectrum of a single WR star ity and progress in modeling are certainly needed exists, the f-i-r emission-line diagnostic that has to understand fully the high-energy processes in been used for O-star winds has not been applied the winds of massive stars at late stages of their to WR stars. X-rays produced close to the stel- evolution. lar core must travel through the bulk of photo- absorbing wind. Tentative evidences for absorp- Based on observations obtained with XMM- tion edges are found in the X-ray spectra of WR1 Newton, an ESA science mission with instruments (Ignace et al. 2003) and WR142. and contributions directly funded by ESA Mem- In the wind of WR 142, the radius at which ber States and NASA. This research has made use the photo-absorption to X-rays is equal to unity is of NASA’s Astrophysics Data System Service and plotted in Fig. 2 against photon energy. The two the SIMBAD database, operated at CDS, Stras- shown models reproduce the stellar SED equally bourg, France. Special thanks to Robert Gru- well. One of the models is thick to X-rays for endl for providing the IR measurements. The au- 1000’s of R∗, and the observed X-rays would have thors thank the referee M. De Becker for impor- to emerge from the far outer regions of the wind. tant comments that improved the paper. Fund- Origin of X-rays far out in the wind is favored ing for this research has been provided by NASA in the alternative scenario of X-ray emission from grant NNX08AW84G (Y-HC and RI) and DLR O stars (Pollock 2007). However, a reduction in grant 50OR0804 (LMO).

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This 2-column preprint was prepared with the AAS LATEX macros v5.2.

7 -12

] 3

-1 -13

o 9.54 A 8.89 2 -2 -14 8.60 Rel. Flux cm 7.63 1 7.23 He II 5 - 4 -1 7.10 -15 0 6.69 10000 10500 o [erg s λ

λ / A -16

log F 5.35 -17

3.5 4.0 4.5 5.0 5.5 o log λ / A

Fig. 1.— Spectral energy distribution of WR 142. Boxes give observed 2MASS, Spitzer IRAC and MIPS magnitudes (labels). Observed optic and IR spectra are shown by thin blue lines. The red line is the continuum flux (plus lines in the UV and op- tical) from a PoWR model with log L/L⊙ =5.35, −5.1 −1 T∗ = 160kK, M˙ = 10 M⊙yr reddened with EB−V = 1.73 mag. The model with log L/L⊙ = −5.4 −1 5.65 and M˙ = 10 M⊙yr gives an equally good fit. A typical He ii line is shown in the inset to illustrate the effect of extreme rotational broad- ening – the observed line (blue) is compared to the synthetic line convolved with v sini = 4000kms−1 (red).

8 2000 Energy [keV] 8 4 2 1 0.5 O IV (K-shell) C IV (K-shell) O IV 2p (b-f) C IV 2s (b-f) C IV 2p (b-f) He II 1s (b-f)

Model boundary

1500 ) 2 o = 1 τ

1000 1 Radius where

500 log (Counts per A

0 0 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 o o log λ/A log(λ/A)]

Fig. 2.— Radius (in units of R∗) where the ra- Fig. 4.— Thick red line: XMM-Newton MOS-2 dial optical depth becomes unity, as a function of spectrum of WR142 where background has not wavelength in the X-ray range. The red curve is been subtracted. Thin blackline: MOS-2 spec- for the same model for which the spectral energy trum of a background region. The error bars give is shown in Fig. 1, while the blue curve is for a 2σ. model with half the mass-loss rate and otherwise similar parameters.

Energy [keV] 0.3 0.5 1 2 3 5 140 120 100 80 60

Counts / keV 40 WR142 20 0 -0.5 0.0 0.5 log (Energy) [keV] Fig. 3.— Part of the merged XMM-Newton EPIC image (0.2-12.0 keV) with over-plotted contours. Fig. 5.— XMM-Newton MOS-2 (red) and PN Image size is 5.7′ × 4.1′. WR142 is marked by (blue) spectra of WR142 with subtracted back- an arrow. The coordinates are equatorial (J2000). ground. The error bars give 2σ North up, east left.

9