arXiv:1202.1525v1 [astro-ph.SR] 7 Feb 2012
High-resolution X-ray spectroscopy reveals the special nature of
Wolf-Rayet star winds
Institute for Physics and Astronomy, University Potsdam, 14476 Potsdam, Germany
Department of Physics and Astronomy, University of Iowa, Iowa City, IA 52245, USA
Institute for Physics and Astronomy, University Potsdam, 14476 Potsdam, Germany
Massachusetts Institute of Technology, Kavli Institute for Astrophysics and Space Research,
70 Vassar St., Cambridge, MA 02139, USA
Department of Physics and Astronomy, East Tennessee State University, Johnson City,
TN 37663, USA
European Space Agency XMM-Newton Science Operations Centre, European Space
Astronomy Centre, Apartado 78, Villanueva de la Ca˜ nada, 28691 Madrid, Spain
We present the first high-resolution X-ray spectrum of a putatively single
Wolf-Rayet star. 400ks observations of WR6 by the XMM-Newton-telescope
resulted in a superb quality high-resolution X-ray spectrum. Spectral analysis
reveals that the X-rays originate far out in the stellar wind, more than 30 stellar
radii from the photosphere, and thus outside the wind acceleration zone where the
line-driving instability could create shocks. The X-ray emitting plasma reaches
temperatures up to 50MK, and is embedded within the un-shocked, “cool” stellar
wind as revealed by characteristic spectral signatures. We detect a fluorescent Fe
– 2 –
line at ≈ 6.4keV. The presence of fluorescence is consistent with a two-component
medium, where the cool wind is permeated with the hot X-ray emitting plasma.
The wind must have a very porous structure to allow the observed amount of
X-rays to escape. We find that neither the line-driving instability nor any alter-
native binary scenario can explain the data. We suggest a scenario where X-rays
are produced when the fast wind rams into slow ”sticky clumps” that resist ac-
celeration. Our new data show that the X-rays in single WR-star are generated
by some special mechanism different from the one operating in the O-star winds.
Subject headings: Stars: winds, outflows — Stars: Wolf-Rayet — Stars: individ-
ual: WR 6 — X-rays: stars
Massive stars reach the Wolf-Rayet (WR) evolutionary phase when their hydrogen fuel
has been consumed or lost and the products of nuclear fusion appear in their atmospheres, be-
fore ending their short lives among the progenitors of core-collapse supernovae (SNe) (Smartt
2009). These stars are millions of times more luminous than the sun and drive strong stel-
lar winds by radiation pressure exerted through absorption and scattering in spectral lines
(Nugis & Lamers 2002; Gr¨ afener & Hamann 2005). This “line-driving” mechanism is known
to be unstable (Lucy & White 1980) and is thought to create hydrodynamic shocks that heat
the plasma to a few million degrees with the emission of thermal X-rays (Feldmeier et al.
1997). For O-type stars, which have weaker winds than WR stars, this model is largely con-
sistent with the observations (Oskinova et al. 2006; G¨ udel & Naz´ e 2009). The winds of WR
stars are also expected to suffer from the line-driving instability (LDI) (Gayley & Owocki
1995) and thus may emit X-rays similar to the O-type stars. This conjecture was not verified
by observations so far.
In this Letter we report the first high-resolution X-ray spectrum of a putatively single
WN-type star and its analysis.
2. The WN star WR6
The object of our study, WR6 (EZCMa, HD50896), has been successfully modeled
as a hydrogen-free WN star of subtype WN5 by fitting its UV, visible, and infrared spec-
tra (Hamann et al. 2006). Its stellar temperature is 90kK, and its luminosity is 105.6L⊙
when adopting a distance of 1.82kpc as implied from its membership to the association
– 3 –
Collinder121. The mass-loss rate is
minal velocity of v∞≈ 1,700kms−1. Hence the wind carries a mechanical power
equivalent to 104L⊙, which is about 3% of the stellar luminosity.
˙M ≈ 2 × 10−5M⊙yr−1, and the wind reaches a ter-
The WR winds are strongly inhomogeneous (L´ epine & Moffat 1999). Narrow spectral
features drifting over the broad emission-line profiles are interpreted as showing radially
accelerated blobs of matter (Moffat et al. 1988). Quasi-periodic variability could be related
to rotation (Dessart & Chesneau 2002) and caused by spiral-like density patterns in the
wind rooted in the photosphere, similar to the corotating interaction regions observed in the
solar wind (Mullan 1984; Morel et al. 1997). WR6 shows considerable photometric, spectral,
and polarimetric variability on the time scale of 3.766d (Duijsens et al. 1996; St-Louis et al.
Its thermal radio spectrum (Dougherty & Williams 2000) and complex variability pat-
tern (Duijsens et al. 1996) support the conclusion that WR6 is not a binary but a single
star. Previous X-ray observations of WR6 at lower spectral resolution confirmed that its
X-ray luminosity and temperature are typical of single WN stars (Pollock 1987; Ignace et al.
2003; Oskinova 2005; Skinner et al. 2010).
3.Observations and Spectral Modeling
The X-ray data on WR6 were taken with the X-Ray Multi-Mirror Satellite XMM-
Newton. Its telescopes illuminate three different instruments which always operate simulta-
neously: RGS is a Reflection Grating Spectrometer, achieving a spectral resolution of 0.07˚ A;
RGS is not sensitive for wavelengths shorter than 5˚ A. The other focal instruments MOS and
PN cover the shorter wavelengths; their spectral resolution is modest (E/∆E ≈20 – 50).
The data were obtained at four epochs in 2010 (Oct.11 and 13, Nov.4 and 6). The total
exposure time of 400 kilosecond was split into four individual parts. Our data reduction
involved standard procedures of the XMM-Newton Science Analysis System v.10.0.
The RGS spectrum is shown in Fig.1. The X-ray luminosity of WR6 in the 0.3 –
12.0keV band is LX ≈ 8 × 1032ergs−1, or 0.1L⊙.
mechanical power. The RGS spectrum is dominated by strong and broad emission lines
of metals in accordance with the WN-wind abundances characteristic of CNO processed
This constitutes 10−5of the wind’s
To model the observed X-ray spectrum we generated emissivities for a hydrogen-free
plasma with the APEC thermal plasma model (Smith et al. 2001). For most elements the
– 4 –
abundances were fixed to the values derived from the UV and optical spectra.
In this stage we approximate the line profiles by Gaussians, with a common Doppler shift
against the laboratory wavelength of −650kms−1and a width (FWHM) of 3000kms−1as
the best values found by the automatic procedure from fitting the RGS data. The optimum
fits to the observed spectra were obtained with the ISIS software (Noble & Nowak 2008).
The observed spectra, the best-fitting model, and the residuals are shown in Fig.2. In
order to achieve an acceptable fit, it was necessary to compose the emission from plasmas of
three different temperatures (1.6, 7.0 and 45MK) with respective emission measures of 253,
86, and 22 ×1054cm−3. The hottest component is required to reproduce the 2–5˚ A continuum
and the emission lines from Fe xxv, Ca xix, Ar xvii, and S xv. Nevertheless, there are still
residuals at some lines, such as Mg xii and Ne x, indicating that a three-temperature model
is not fully sufficient, or that some abundances are not well determined. O and C lines are
absent or very weak, as expected for a plasma with the chemical composition of a WN-type
To account for the absorption within the wind we applied the vphabs model (Arnaud
1996), modified for the absence of hydrogen. Strong absorption, significantly exceeding that
of the interstellar medium towards the star, is evident. The ratio of fluxes in the Nvi and
Nvii Lyα lines is reproduced when absorption in the cool wind is included, especially the
K-shell ionization edge of Niv that is located between the two aforementioned lines (see
The low-resolution EPIC spectra provide data up to 12keV. Significant residuals are
encountered in the region around the Fe xxv line at about 1.9˚ A (see Fig.2). The observation
shows a definitely broader feature than reproduced by the plasma model. Obviously, this
complex includes a further blend. The fit of an additional gaussian profile yields a flux
of 4.5 × 10−7photonscm−2s−1and a line center at 1.92˚ A. The latter value agrees within
the uncertainties with the wavelength of the Fe Kα line (1.94˚ A). Hence we conclude that
cool-wind material must be irradiated by X-rays which are hard enough (> 7.1keV) to
excite this line by fluorescence. Our PoWR wind models show that in the outer wind the
leading ionization stage of iron is Fev. The presence of fluorescence is consistent with a
two-component medium, where the cool wind is permeated with the hot X-ray emitting
– 5 –
4. The cool wind model of WR6
To determine the physical conditions in the “cool” component of the stellar wind, we
employ the PoWR model atmosphere code (Hamann & Gr¨ afener 2004). The PoWR code
solves the non-LTE radiative transfer simultaneously with the equations of statistical and
radiative equilibrium. Complex model atoms with thousands of transitions are taken into
account. The extensive inclusion of the iron group elements is important because of their
blanketing effect on the atmospheric structure (Gr¨ afener et al. 2002). A particular stellar
atmosphere model is defined by the effective temperature, surface gravity, luminosity, mass-
loss rate, wind terminal velocity, and chemical composition.
For the supersonic part of the wind, we adopt the radial dependence of velocity as
v(r) = v∞(1 − 1/r), where r is in units of the stellar radius R∗, and the terminal velocity
v∞is a free parameter. In the subsonic region, the velocity field is defined such that the
hydrostatic density stratification is approached. The PoWR models account for stellar-wind
clumping in the standard volume-filling factor ‘microclumping’ approximation (Hillier 1991;
Hamann & Koesterke 1998). Our best model of WR6 has a clumping factor D = 20, where
D = ?ρ2?/?ρ?2. The synthetic spectra are calculated over the whole spectral range from UV
to IR, and then compared to the observed spectra and photometric fluxes.
With the final wind model being adopted, the stratifications of the density, opacity,
ionization, and radiative flux are defined. The mass absorption increases with distance from
the star, partly because the fully ionized helium recombines to He ii with its strong bound-
free opacity. This increase of the mass absorption coefficient partly compensates for the
decrease of matter density, keeping the wind opaque to large radii.
In Fig.3 (upper panel) we show the radius in the wind where the continuum optical
depth on the radial ray becomes unity, as predicted by PoWR model. Roughly speaking,
only X-rays emitted from outside that radius can escape from the wind and be seen by a
distant observer, unless inhomogeneities create some porosity (“macroclumping”).
5. The analysis of line ratios in He-like ions
Helium-like ions show a group of three lines, consisting of a forbidden (f), an intercom-
bination (i), and a resonance (r) transition – the so-called fir triplet. The fir triplets of
Si xiii, Mg xi, Ne ix, and N vi are present in the RGS spectrum of WR6. As illustrated
in Fig.4, the line ratios observed in WR6 are very different from those observed in typical
single O-type stars. For Nvi the ratio f/i exceeds unity in WR6, while O stars show f/i < 1
(Leutenegger et al. 2006; Waldron & Cassinelli 2007).
– 6 –
For each He-like ion, the ratio of fluxes between the forbidden and the intercombination
component, R, is sensitive to the electron density and the ultraviolet flux (Blumenthal et al.
1 + φ(r)/φc+ Ne(r)/Nc,(1)
where φ is the photoexcitation rate from the term 2s3S to 2p3P, and Ne is the electron
density. The quantities R0,φc, and Ncdepend only on atomic parameters and the electron
temperature (Blumenthal et al. 1972; Porquet et al. 2001).
We use the PoWR model to calculate the values of R(r) for Mg xi, Ne ix, and N vi
as function of the radial location in the wind of WR6. The density in the X-ray emitting
shocks is not known, therefore we computed two sets of R(r), one neglecting the density
term, and another one assuming that the density in the shock is the same as in the ambient
“cool” wind but the plasma is fully ionized. According to the strong-shock condition, the
hot gas density would be even a factor of four higher than that (Zeldovich & Raizer 1966).
The photo-excitation rates φ(r) are computed at each radius from the radiation intensity as
provided by the PoWR model, which accounts not only simply for geometric dilution, but
also for the diffuse radiative field. Fig.3 (low panel) shows for the Nvi triplet the predicted
R ratio as function of the radial location of the emitting plasma with the best-fit value of
≈ 200R∗when the collisions are accounted for. The Nvi emission line indicates a plasma
temperature of ≈ 1.58MK, which could be reached in a strong shock with a velocity jump
by ≈ 400kms−1.
6.Modeling of X-ray emission line profiles
The broadening of the X-ray lines from WR6 is consistent with the known terminal
velocity of the wind of ≈ 1700kms−1. The lines centers are displaced from their laboratory
wavelengths by ≈ −0.06˚ A (see Fig.5). Such blue-shifts are predicted if the line emission
is distributed within a partly absorbing wind (e.g. Macfarlane et al. 1991; Ignace 2001).
However, according to the standard model, the wind of WR6 remains optically thick till
large distances from the star (550 – 900R∗for λ > 20˚ A). X-rays can hardly emerge from
interior to this location unless the wind is very porous.
We employed our 2-D stochastic wind code (Oskinova et al. 2004) to model the observed
X-ray lines. Relaxing the microclumping approximation, the clumps may have arbitrary
optical depth (“macroclumping”). This wind fragmentation alters the radiative transfer
drastically, compared to a homogeneous wind with the same mass-loss rate (Feldmeier et al.
2003). The X-rays can escape from deeper in the wind than indicated in the upper panel in
– 7 –
Fig.3 for the smooth-wind model.
We assume that parcels of X-ray emitting gas are distributed between clumps of continuum-
absorbing, cool wind. The wind is fully fragmented throughout the whole relevant radial
range. The X-ray emission is placed only at radii larger than rem. Since thermal X-ray
emission arises from the decay of collisionally excited levels and from free-free emission, is
plausible to assume that the emissivity scales as the square of the density, ηλ∝ ρ2. With this
scaling, most X-rays are released close to rem. The intensity along any ray is reduced by the
clumps which are encountered.
λ), where τc
λis the summation over the optical depths through all individual
A grid of line models was calculated for different values of the radius remand ?Nc? –
the time-averaged number of fragments in a radial direction. The radial dependence of the
mass-absorption coefficient and the wind opacity was accounted for in the models. Even with
porosity, emission arising from too close to the star is totally trapped, while the adopted
scaling of the emissivity with ρ2requires an unrealistic amount of energy being converted
into X-rays. This excludes small values of rem. In the other extreme, when X-ray emission
starts only at radii where the wind is already nearly transparent, the observed blue-shift and
asymmetry of the line profiles cannot be reproduced. The model line profiles shown in Fig.5
were calculated with rem= 30R∗and ?Nc? = 50.
7.On The Origin of X-ray Emission in WR stars
Observed temperatures up to 50MK require shock velocity jumps comparable to the
wind speed. From UV spectra of WR6 there is direct evidence that a fraction of the
wind flows with much higher velocity (up to 3,100kms−1) than the general wind speed
(1,700kms−1) (Prinja et al. 1990). We suggest that the hottest gas is produced when very
fast flows impacts on slowly moving clumps within the wind acceleration zone. Such extreme
shocks cannot be maintained to large radii, but the hardest X-rays can escape from the inner
wind owing to smaller absorption cross-sections at higher energies.
Exceptionally large, pancake-shaped clumps which form close to the photosphere at
low velocities would resist the radiative acceleration. Being an obstacle to the generally
much faster wind, they create strong reverse shocks that may add more mass to the clump
if the gas cools quickly, or generate bow shocks if it does not (Cassinelli et al. 2008). Self-
consistent kinematic models of Gayley (in prep.) show that 500kms−1shocks may be present
at ∼ 50R∗in the wind, suitable for X-ray escape if the wind is sufficiently porous.
– 8 –
8. On the binary hypothesis for WR6
The X-ray emission of WR6 varies considerably (Pollock 1987; Willis & Stevens 1996).
Our new observations confirm and extend these results. The XMM-Newton light curve shows
different variability time scales, including a variation of 20% over a month’s span. We also
observe some spectral changes. The X-ray variability does not follow the 3.766d period
which is seen in the optical and UV (Willis et al. 1989). The X-ray variability of WR6
further highlights the differences between single WR and O stars. The latter are remarkably
constant X-ray sources (Naz´ e et al. 2012, submitted).
The optical variability of WR6 has long been known (Wilson 1948), and it appears to
be incompatible with stellar pulsations or the Be star phenomenology (Robert et al. 1992).
Most plausibly, the variability is due to changes of wind structures (St.-Louis et al. 1993;
Flores et al. 2011). This explanation is in line with our new X-ray observations.
Firmani et al. (1980) established that if the 3.766day period were associated with bi-
narity, the mass of the secondary star could not be larger than ≈ 1.5M⊙. A star of such
mass and age (assuming coeval formation of both binary components) would be either a
degenerate neutron star (NS) or a pre main-sequence young stellar object (TTauri type
The NS scenario was criticized because of the relatively low X-ray luminosity of WR6
and the lack of evidence in its UV spectrum (Stevens & Willis 1988). The WR+TTau
scenario is also extremely unlikely, as the high-resolution X-ray spectrum of WR6 is incom-
patible with the X-ray properties of young stellar objects (Kastner et al. 2002).
Variability is a common property of WR stars. Among rigorously monitored WN stars,
40% show optical variability similar to WR6 (Chen´ e & St-Louis 2011). The assumption that
40% of all WN stars have a NS or a TTauri type companion is not realistic.
Our high-resolution X-ray spectra show that the X-ray emitting plasma moves at about
the same velocity as the cool wind. The X-ray line blue-shifts do apparently not change with
time. The X-ray emission line spectrum is compatible with the WN star abundances. All
these facts add further evidence against WR6 being a binary system.
Assuming that WR6 is typical, our results show that in the winds of WN stars strong
shocks are active at large radii from stellar photosphere. A possible explanation involves slow,
dense clumps that accumulate from the wind close above the photosphere without being ef-
– 9 –
fectively accelerated. The seeding of such slow clumps might be associated with (sub)surface
convective zones (Cantiello & Braithwaite 2011). The effects of such wind anisotropies on
the distribution of circumstellar matter could influence the behavior of SNIbc.
In any case, from high-resolution X-ray spectral data of WR6 we conclude that unknown
mechanisms must operate in WR winds. Identifying these mechanisms poses a challenging
problem for the theory of stellar winds.
Based on observations made with XMM-Newton, an ESA science mission with instru-
ments and contributions directly funded by ESA member states and the USA (NASA). We
thank R.K. Smith and A.Foster for help in computing emissivities for a hydrogen-depleted
plasma. DPH was supported by NASA through the Smithsonian Astrophysical Observatory
contract SV3-73016 for the Chandra X-Ray Center and Science Instruments. Funding for
this research has been provided by DLR grant 50OR1101 (LMO).
Arnaud, K. A. 1996, in Astronomical Society of the Pacific Conference Series, Vol. 101,
Astronomical Data Analysis Software and Systems V, ed. G. H. Jacoby & J. Barnes,
Blumenthal, G. R., Drake, G. W. F., & Tucker, W. H. 1972, ApJ, 172, 205
Cantiello, M., & Braithwaite, J. 2011, A&A, 534, A140
Cassinelli, J. P., Ignace, R., Waldron, W. L., et al. 2008, ApJ, 683, 1052
Chen´ e, A.-N., & St-Louis, N. 2011, ApJ, 736, 140
Dessart, L., & Chesneau, O. 2002, A&A, 395, 209
Dougherty, S. M., & Williams, P. M. 2000, MNRAS, 319, 1005
Duijsens, M. F. J., van der Hucht, K. A., van Genderen, A. M., et al. 1996, A&AS, 119, 37
Feldmeier, A., Oskinova, L., & Hamann, W.-R. 2003, A&A, 403, 217
Feldmeier, A., Puls, J., & Pauldrach, A. W. A. 1997, A&A, 322, 878
Firmani, C., Koenigsberger, G., Bisiacchi, G. F., Moffat, A. F. J., & Isserstedt, J. 1980,
ApJ, 239, 607
– 10 –
Flores, A., Koenigsberger, G., Cardona, O., & de La Cruz, L. 2011, Rev. Mexicana Astron.
Astrofis., 47, 261
Gayley, K. G., & Owocki, S. P. 1995, ApJ, 446, 801
Gr¨ afener, G., & Hamann, W.-R. 2005, A&A, 432, 633
Gr¨ afener, G., Koesterke, L., & Hamann, W.-R. 2002, A&A, 387, 244
G¨ udel, M., & Naz´ e, Y. 2009, A&A Rev., 17, 309
Hamann, W.-R., & Gr¨ afener, G. 2004, A&A, 427, 697
Hamann, W.-R., Gr¨ afener, G., & Liermann, A. 2006, A&A, 457, 1015
Hamann, W.-R., & Koesterke, L. 1998, A&A, 335, 1003
Hillier, D. J. 1991, A&A, 247, 455
Ignace, R. 2001, ApJ, 549, L119
Ignace, R., Oskinova, L. M., & Brown, J. C. 2003, A&A, 408, 353
Kastner, J. H., Huenemoerder, D. P., Schulz, N. S., Canizares, C. R., & Weintraub, D. A.
2002, ApJ, 567, 434
L´ epine, S., & Moffat, A. F. J. 1999, ApJ, 514, 909
Leutenegger, M. A., Paerels, F. B. S., Kahn, S. M., & Cohen, D. H. 2006, ApJ, 650, 1096
Lucy, L. B., & White, R. L. 1980, ApJ, 241, 300
Macfarlane, J. J., Cassinelli, J. P., Welsh, B. Y., et al. 1991, ApJ, 380, 564
Moffat, A. F. J., Drissen, L., Lamontagne, R., & Robert, C. 1988, ApJ, 334, 1038
Morel, T., St-Louis, N., & Marchenko, S. V. 1997, ApJ, 482, 470
Mullan, D. J. 1984, ApJ, 283, 303
Noble, M. S., & Nowak, M. A. 2008, PASP, 120, 821
Nugis, T., & Lamers, H. J. G. L. M. 2002, A&A, 389, 162
Oskinova, L. M. 2005, MNRAS, 361, 679
Oskinova, L. M., Feldmeier, A., & Hamann, W.-R. 2004, A&A, 422, 675
– 11 –
—. 2006, MNRAS, 372, 313
Pollock, A. M. T. 1987, ApJ, 320, 283
Porquet, D., Mewe, R., Dubau, J., Raassen, A. J. J., & Kaastra, J. S. 2001, A&A, 376, 1113
Prinja, R. K., Barlow, M. J., & Howarth, I. D. 1990, ApJ, 361, 607
Robert, C., Moffat, A. F. J., Drissen, L., et al. 1992, ApJ, 397, 277
Skinner, S. L., Zhekov, S. A., G¨ udel, M., Schmutz, W., & Sokal, K. R. 2010, AJ, 139, 825
Smartt, S. J. 2009, ARA&A, 47, 63
Smith, R. K., Brickhouse, N. S., Liedahl, D. A., & Raymond, J. C. 2001, ApJ, 556, L91
St-Louis, N., Chen´ e, A.-N., Schnurr, O., & Nicol, M.-H. 2009, ApJ, 698, 1951
St.-Louis, N., Howarth, I. D., Willis, A. J., et al. 1993, A&A, 267, 447
Stevens, I. R., & Willis, A. J. 1988, MNRAS, 234, 783
Waldron, W. L., & Cassinelli, J. P. 2007, ApJ, 668, 456
Willis, A. J., Howarth, I. D., Smith, L. J., Garmany, C. D., & Conti, P. S. 1989, A&AS, 77,
Willis, A. J., & Stevens, I. R. 1996, A&A, 310, 577
Wilson, O. C. 1948, PASP, 60, 383
Zeldovich, Y. B., & Raizer, Y. P. 1966, Elements of gasdynamics and the classical theory of
shock waves, ed. Zeldovich, Y. B. & Raizer, Y. P.
This preprint was prepared with the AAS LATEX macros v5.2.
– 12 –
N VIIN VII
Fig. 1.— The RGS spectrum of WR6 with the strongest emission lines identified. The error
bars correspond 3σ. The RGS1 spectrum is shown from 18.2˚ A to 26.7˚ A, otherwise the
RGS2 spectrum is plotted
– 13 –
Fig. 2.— The X-ray spectrum of WR6, obtained with different detectors: RGS (top panel),
MOS (middle), and PN (bottom). Black curves display the observations; the model is shown
in red. The lower part of each panel displays the residuals. The PN panel includes an inset to
the right showing the transmission of the cool wind plus interstellar absorption for λ > 20˚ A.
The N iv K-shell absorption edge is located between the N vi and N vii lines (red vertical
bars). An insert to the left in the PN panel shows a broad spectral feature that is not
reproduced by the thermal plasma model (blue dotted line). An additional line at 1.92˚ A,
which we identity as fluorescent FeKα, improves the fit (red histogram).
– 14 –
K-SHELL N V
K-SHELL N IV
K-SHELL N III
Radius where τ = 1 [R*]
R = 4.81
Ne = 0
Ne as in cool wind, fully ionized
log radius [R*]
R (line ratio f/i)
Fig. 3.— Upper panel: radius where the continuum optical depth reaches unity, as predicted
by the PoWR model for the “cool” wind component of WR6, when inhomogeneities are only
accounted for in the “microclumping” approximation. Lower panel: Dependence of the line
ratio R = f/i for the Nvi lines as a function of the radial location of the emitting plasma.
Based on the PoWR model, the red curve accounts only for radiative de-population, while
the blue curve includes also collisional process under the assumption that the shocked plasma
has the same density as the cool wind. The measured value is shown as horizontal green
line, with the green shaded area representing the 3σ confidence band of the measurement.
– 15 –
28.4 28.6 28.8 29.0 29.2 29.4 29.6
Fig. 4.— Upper panel: RGS spectrum of WR6 in the range of the Nvi triplet of lines
(black). The vertical error bars correspond to 3σ. The vertical lines indicate the restframe
wavelength of the resonance (λr), intercombination (λi), and forbidden (λf) line, respectively.
The solid red line shows a fit to these lines by a formal model with three Gaussian profiles.
Lower panel: Same as upper, but for the O4I star ζ Puppis. The red line shows the scaled
Gaussian fit from the upper panel.