arXiv:astro-ph/0610356v2 21 Feb 2007 Wolf-Rayet Stars 1 Physical Properties of Wolf-Rayet Stars Paul A. Crowther Department of Physics & Astronomy, University of Sheffield, Hounsfield Road, Sheffield, S3 7RH, United Kingdom, email: Paul.Crowther@sheffield.ac.uk Key Words stars: Wolf-Rayet stars: fundamental parameters stars: evolu- tion stars: abundances Abstract The striking broad emission line spectroscopic appearance of Wolf-Rayet (WR) stars has long defied analysis, due to the extreme physical conditions within their line and continuum forming regions. Recently, model atmosphere studies have advanced sufficiently to enable the determination of stellar temperatures, luminosities, abundances, ionizing fluxes and wind properties. The observed distributions of nitrogen (WN) and carbon (WC) sequence WR stars in the Milky Way and in nearby star forming galaxies are discussed these imply lower limits to progenitor masses of ���25, 40, 75 M��� for hydrogen-depleted (He-burning) WN, WC, and H-rich (H-burning) WN stars, respectively. WR stars in massive star binaries permit studies of wind-wind interactions and dust formation in WC systems. They also show that WR stars have typical masses of 10���25M���, extending up to 80M��� for H-rich WN stars. Theoretical and observational evidence that WR winds depend on metallicity is presented, with implications for evolutionary models, ionizing fluxes, and the role of WR stars within the context of core-collapse supernovae and long-duration gamma ray bursts. CONTENTS Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 3

2 Paul A. Crowther Observed Properties . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5 Spectral Properties and Spectral Classification . . . . . . . . . . . . . . . . . . . . . . . 5 Absolute magnitudes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7 Observed distribution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8 Binary statistics and masses . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 14 Rotation velocities . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15 Stellar wind bubbles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15 Physical Parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16 Radiative Transfer . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16 Stellar Temperatures and Radii . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17 Stellar Luminosities . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20 Ionizing fluxes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21 Elemental abundances . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23 Wind Properties . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27 Wind velocities . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27 Mass-loss rates . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29 Clumping . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 30 Metallicity dependent winds? . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31 Line driving in WR winds . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35 Interacting Binaries . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39 Close binary evolution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39 Colliding winds . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40 Dust formation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 42 Evolutionary models and properties at core-collapse . . . . . . . . . . . . . . . . . 44 Rotational mixing . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44 Evolutionary model predictions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46 WR stars as SNe and GRB progenitors . . . . . . . . . . . . . . . . . . . . . . . . . . 48

Annu. Rev. Astro. Astrophys. Sept 2007 Vol 45 Summary Points . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51 Future Issues to be Resolved . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 52 1 Introduction Massive stars dominate the feedback to the local interstellar medium (ISM) in star-forming galaxies via their stellar winds and ultimate death as core-collapse supernovae. In particular, Wolf-Rayet (WR) stars typically have wind densities an order of magnitude higher than massive O stars. They contribute to the chemical enrichment of galaxies, they are the prime candidates for the immediate progenitors of long, soft Gamma Ray Bursts (GRBs, Woosley & Bloom 2006), and they provide a signature of high-mass star formation in galaxies (Schaerer & Vacca 1998). Spectroscopically, WR stars are spectacular in appearance, with strong, broad emission lines instead of the narrow absorption lines which are typical of normal stellar populations (e.g. Beals 1940). The class are named after Wolf & Rayet (1867) who identified three stars in Cygnus with such broad emission lines. It was immediately apparent that their spectra came in two flavours, subsequently identified as those with strong lines of helium and nitrogen (WN subtypes) and those with strong helium, carbon, and oxygen (WC and WO subtypes). Gamov (1943) first suggested that the anomalous composition of WR stars was the result of nuclear processed material being visible on their surfaces, although this was not universally established until the final decade of the 20th Century (Lamers et al. 1991). Specifically, WN and WC stars show the products of the CNO cycle (H-burning) and the triple-�� (He-burning), respectively. In reality, there is a continuity of physical and chemical properties between O supergiants and WN 3

4 Paul A. Crowther subtypes. Typically, WR stars have masses of 10���25 M���, and are descended from O-type stars. They spend ���10% of their ���5Myr lifetime as WR stars (Meynet & Maeder 2005). At Solar metallicity the minimum initial mass for a star to become a WR star is ���25 M���. This corresponds closely to the Humphreys & Davidson (1979) limit for red supergiants (RSG), according to a comparison between the current temperature calibration of RSG and stellar models that allow for mass-loss and rotation (e.g. Levesque et al. 2005). Consequently, some single WR stars are post-red supergiants within a fairly limited mass range of probably 25���30M���. Evolution proceeds via an intermediate Luminous Blue Variable (LBV) phase above 30M���. For close binaries, the critical mass for production of a WR star has no such robust lower limit, since Roche lobe overflow or common envelope evolution could produce a WR star instead of an extended RSG phase. The strong, broad emission lines seen in spectra of WR stars are due to their powerful stellar winds. The wind is su���ciently dense that an optical depth of unity in the continuum arises in the outflowing material. The spectral features are formed far out in the wind and are seen primarily in emission. The line and continuum formation regions are geometrically extended compared to the stellar radii and their physical depths are highly wavelength dependent. The unique spectroscopic signature of WR stars has permitted their detection individually in Local Group galaxies (e.g. Massey & Johnson 1998 Massey 2003), collectively within knots of local star forming galaxies (e.g. Hadfield & Crowther 2006), and as significant contributors to the average rest-frame UV spectrum of Lyman Break Galaxies (Shapley et al. 2003). The present review focuses on observational properties of classical Wolf-Rayet

Wolf-Rayet Stars 5 stars in the Milky Way and beyond, plus physical and chemical properties deter- mined from spectroscopic analysis, plus comparisons with interior evolutionary models, and provides revisions to the topic with respect to the excellent Abbott & Conti (1987) review. Low mass (��� 0.6M���) central stars of Planetary Nebulae displaying a Wolf-Rayet spectroscopic appearance (denoted [WR]) are not con- sidered. Nevertheless, analysis tools discussed here are common to both types of star (e.g. Crowther et al. 2006a). 2 Observed Properties 2.1 Spectral Properties and Spectral Classification Visual spectral classification of WR stars is based on emission line strengths and line ratios following Smith (1968a). WN spectral subtypes follow a scheme involving line ratios of N iii-v and He i-ii, ranging from WN2 to WN5 for ���early WN��� (WNE) stars, and WN7 to WN9 for ���late WN��� (WNL) stars, with WN6 stars either early or late-type. A ���h��� su���x may be used to indicate the presence of emission lines due to hydrogen (Smith, Shara & Moffat 1996). Complications arise for WN stars with intrinsically weak emission lines. For example, WR24 (WN6ha) has a He ii ��4686 emission equivalent width that is an order of magnitude smaller than those in some other WN6 stars the ���ha��� nomen- clature indicates that hydrogen is seen both in absorption and emission. From a standard spectroscopic viewpoint, such stars possess mid to late WN spec- tral classifications. However, their appearance is rather more reminiscent of Of stars than classic WN stars, since there exists a continuity of properties between normal O stars and late-type WN stars. These stars are widely believed to be massive O stars with relatively strong stellar winds at a rather early evolutionary

6 Paul A. Crowther stage. They are believed not to represent the more mature, classic He-burning WN stars. Smith, Crowther & Prinja (1994) extended the WN sequence to very late WN10���11 subtypes in order to include a group of emission line stars originally classified as Ofpe/WN9 (Bohannan & Walborn 1989). WN11 subtypes closely resemble extreme early-type B supergiants except for the presence of He ii ��4686 emission. A quantitative comparison of optical line strengths in Of and WNL stars is presented in figure 8 of Bohannan & Crowther (1999). R127 (WN11) in the Large Magellanic Cloud (LMC) was later identified as a LBV (Stahl et al. 1983), whilst a famous Galactic LBV, AG Car exhibited a WN11-type spectrum at visual minimum (Walborn 1990 Smith et al. 1994). Various multi-dimensional classification systems have been proposed for WN stars they generally involve line strengths or widths, such that strong/broad lined stars have been labelled WN-B (Hiltner & Schild 1966), WN-s (Hamann, Koesterke & Wessolowski 1993) or WNb (Smith, Shara & Moffat 1996). Of these, none have generally been adopted. From a physical perspective, strong- and weak-lined WN stars do form useful sub-divisions. Therefore we shall define weak (-w) and strong (-s) WN stars as those with He ii ��5412 equivalent widths smaller than or larger than 40 �� A. An obvious limitation of such an approach is that intrinsically strong-lined WN stars would be diluted by binary companions or nearby stars in spatially crowded regions and so might not be identified as such. WNE-w stars tend to exhibit triangular line profiles rather than the more typical Gaussian lines of WNE-s stars (Marchenko et al. 2004), since one observes material much closer to the stellar core that is being strongly accelerated. WC spectral subtypes depend on the line ratios of C iii and C iv lines along

Wolf-Rayet Stars 7 with the appearance of O iii-v, spanning WC4 to WC9 subtypes, for which WC4��� 6 stars are ���early��� (WCE) and WC7���9 are ���late��� (WCL). Rare, oxygen-rich WO stars form an extension of the WCE sequence, exhibiting strong O vi ����3811-34 emission (Kingsburgh, Barlow & Storey 1995). The most recent scheme involves WO1 to WO4 subtypes depending on the relative strength of O v-vi and C iv emission lines (Crowther, De Marco & Barlow 1998). Finally, C iv ��5801-12 appears unusually strong in an otherwise normal WN star in a few cases, leading to an intermediate WN/C classification (Conti & Massey 1989). WN/C stars are indeed considered to be at an intermediate evolutionary phase between the WN and WC stages. Representative examples of WN and WC stars are presented in Figure 1. Var- ious X-ray to mid-IR spectroscopic datasets of Galactic Wolf-Rayet stars are presented in Table 1, including extreme ultraviolet synthetic spectra from model atmospheres (Smith, Norris & Crowther 2002 Hamann & Gr��afener 2004). 2.2 Absolute magnitudes WR stars cannot be distinguished from normal hot stars using UBV photometry. Broad-band visual measurements overestimate the true continuum level in ex- treme cases by up to 1 magnitude, or more typically 0.5 mag for single early-type WR stars due to their strong emission-line spectra. Consequently, Westerlund (1966) introduced narrow-band ubyr filters that were specifically designed to minimize the effect of WR emission lines (although their effect cannot be entirely eliminated). These passbands were later refined by Smith (1968b) and by Massey (1984), such that most photometry of WR stars has used the ubvr filter system, which is compared to Johnson UBV filters in Fig. 1.

8 Paul A. Crowther As with normal stars, ubv photometry permits a determination of the inter- stellar extinction, Av. Let us adopt a typical ratio of total, AV to selective, E(B ��� V ) = AB ��� AV extinction, RV = AV /E(B ��� V ) = 3.1. Following Turner (1982), the broad-band and narrow-band optical indices for WR stars are then related by: Av = 4.12 Eb���v = 3.40 EB���V = 1.11 AV A direct determination of WR distances via stellar parallax is only possible for �� Vel (WC8+O) using Hipparcos, and even that remains controversial (Millour et al. 2007). Otherwise, cluster or association membership is used to provide an approximate absolute magnitude-spectral type calibration for Milky Way WR stars. The situation is much better for WR stars in the Magellanic Clouds, al- though not all subtypes are represented. Typical absolute magnitudes range from Mv = ���3 mag at earlier subtypes to ���6 mag for late subtypes, or exceptionally ���7 mag for hydrogen-rich WN stars. The typical spread is ��0.5 mag at individual subtypes. 2.3 Observed distribution Conti (1976) first proposed that a massive O star may lose a significant amount of mass via stellar winds, revealing first the H-burning products at its surface, and subsequently the He-burning products. These evolutionary stages are spectro- scopically identified with the WN and WC types. This general picture has since become known as the ���Conti scenario���. Such stars should be over-luminous for their mass, in accord with observations of WR stars in binary systems. Massey (2003) provides a more general overview of massive stars within Local Group galaxies.

Wolf-Rayet Stars 9 2.3.1 WR stars in Milky Way Wolf-Rayet stars are located in or close to massive star forming regions within the Galactic disk. A catalogue is provided by van der Hucht (2001). A quarter of the known WR stars in the Milky Way reside within massive clusters at the Galactic centre or in Westerlund 1 (van der Hucht 2006). From membership of WR stars in open clusters, Schild & Maeder (1984) and Massey, DeGioia-Eastwood & Waterhouse (2001) investigated the initial masses of WR stars empirically. A revised compilation is provided in Crowther et al. (2006b). Overall, hydrogen-rich WN stars (WNha) are observed in young, massive clusters their main-sequence turn-off masses (based on Meynet et al. 1994 isochrones) suggest initial masses of 65 ��� 110M���, and are believed to be core-H burning (Langer et al. 1994 Crowther et al. 1995a). Lower-mass progenitors of 40���50M��� are suggested for classic mid-WN, late WC, and WO stars. Progenitors of some early WN stars appear to be less massive still, suggesting an initial-mass cutoff for WR stars at Solar metallicity around 25M���. From an evolutionary perspective, the absence of RSGs at high luminosity and presence of H-rich WN stars in young massive clusters suggests the following variation of the Conti scenario in the Milky Way, i.e. for stars initially more massive than ��� 75M��� O ��� WN(H ��� rich) ��� LBV ��� WN(H ��� poor) ��� WC ��� SNIc, whereas for stars of initial mass from ��� 40 ��� 75M���, O ��� LBV ��� WN(H ��� poor) ��� WC ��� SNIc, and for stars of initial mass in the range 25���40M���, O ��� LBV/RSG ��� WN(H ��� poor) ��� SNIb.

10 Paul A. Crowther Indeed, the role of the LBV phase is not yet settled ��� it may be circumvented entirely in some cases it may follow the RSG stage, or it may even dominate pre- WR mass-loss for the most massive stars (Langer et al. 1994 Smith & Owocki 2006). Conversely, the presence of dense, circumstellar shells around Type IIn SN indicates that some massive stars may even undergo core-collapse during the LBV phase (Smith et al. 2007). Remarkably few Milky Way clusters host both RSG and WR stars, with the notable exception of Westerlund 1 (Clark et al. 2005) this suggests that the mass range common to both populations is fairly narrow. Although optical narrow-band surveys (see below) have proved very successful for identifying WR stars in the Solar neighbourhood, only a few hundred WR stars are known in the Milky Way, whilst many thousands are expected within the Galactic disk (van der Hucht 2001). Consequently, near-IR narrow-band imaging surveys together with spectroscopic follow-up may be considered for more extensive surveys to circumvent high interstellar extinction (Homeier et al. 2003). Limitations of IR emission-line surveys are that fluxes of near-IR lines are much weaker than those of optical lines, Also, no strong WR lines are common to all spectral types in the frequently used K band. An added complication is that some WC stars form dust which heavily dilutes emission line fluxes longward of the visual. Nevertheless, infrared surveys are presently underway to get an improved census of WR stars in the Milky Way. Alternatively, WR candidates may be identified from their near- to mid-IR colours, which, as in other early-type supergiants, are unusual due to strong free-free excess emission (Hadfield et al. 2007).

Wolf-Rayet Stars 11 2.3.2 WR stars in the Local Group WR stars have typically been discovered via techniques sensitive to their unusually broad emission-line spectra, based on objective prism searches or interference filter imaging (see Massey 2003). Narrow-band interference filter techniques have been developed (e.g. Moffat, Seggewiss & Shara 1985 Massey, Armandroff & Conti 1986) that distinguish strong WR emission lines at He ii ��4686 (WN stars) and C iii ��4650 (WC stars) from the nearby continuum. Such techniques have been applied to regions of the Milky Way disk, the Magellanic Clouds and other nearby galaxies. An example of this approach is presented in Figure 2 for the spiral galaxy NGC 300 (d ��� 2 Mpc). A wide-field image of NGC 300 is presented, with OB complex IV-V indicated, together with narrow-band images centred at ��4684 (He ii 4686) and ��4781 (continuum). Several WR stars are seen in the difference (He ii-continuum) image, including an apparently single WC4 star (Schild et al. 2003). It is well established that the absolute number of WR stars and their sub- type distribution are metallicity dependent. N(WR)/N(O)���0.15 in the relatively metal-rich Solar Neighbourhood (Conti et al. 1983 van der Hucht 2001), yet N(WR)/N(O)���0.01 in the metal-deficient SMC on the basis of only 12 WR stars (Massey, Olsen & Parker 2003) versus ���1000 O stars (Evans et al. 2004). It is believed that the majority of Galactic WR stars are the result of single-star evolution, yet some stars (e.g. V444 Cyg) result from close binary evolution (Vanbeveren et al. 1998). Similar relative numbers of WN to WC stars are observed in the Solar Neigh- bourhood (Hadfield et al. 2007). In contrast, WN stars exceed WC stars by a factor of ���5 and ���10 for the LMC and SMC, respectively (Breysacher, Azzopardi & Testor 1999 Massey, Olsen & Parker 2001). At low metallicity the reduced WR

12 Paul A. Crowther population and the relative dominance of WN subtypes most likely result from the metallicity dependence of winds from their evolutionary precursors (Mokiem et al. 2007). Consequently, only the most massive single stars reach the WR phase in metal-poor environments. Single stars reaching the WC phase at high metallicity may end their lives as a RSG or WN stars in a lower metallicity en- vironment. As such, one might suspect that most WR stars at low metallicity are formed via binary evolution. However, Foellmi, Moffat & Guerrero (2003a) suggest a similar WR binary fraction for the SMC and Milky Way. Not all WR subtypes are observed in all environments. Early WN and WC subtypes are preferred in metal-poor galaxies, such as the SMC (Massey et al. 2003), while late WC stars are more common at super-Solar metallicities, such as M83 (Hadfield et al. 2005) Line widths of early WC and WO stars are higher than late WC stars, although width alone is not a defining criterion for each spectral type. The correlation between WC subclass and line width is nevertheless strong (Torres, Conti & Massey 1986). The subtype distributions of WR stars in the Solar Neighbourhood, LMC, and SMC are presented in Figure 3. We shall address this aspect in Sect 4.4. 2.3.3 WR galaxies Individual WR stars may, in general, be resolved in Local Group galaxies from ground-based observations, whilst the likelihood of contamination by nearby sources increases at larger distances. For example, a typical slit width of 1������ at the 2Mpc distance of NGC 300 corresponds to a spatial scale of ���10 pc. Relatively isolated WR stars have been identified, albeit in the minority (recall Figure 2). This is even more problematic for more distant galaxies such as M 83 where the great majority of WR stars are observed in clusters or associations (Hadfield et al. 2005). So-called ���WR galaxies��� are typically starburst

Wolf-Rayet Stars 13 regions exhibiting spectral features from tens, hundreds, or even thousands of WR stars (Schaerer, Contini & Pindao 1999). Average Milky Way/LMC WN or WC line fluxes (Schaerer & Vacca 1998) are typically used to calculate stellar populations in WR galaxies. These should be valid provided that the line fluxes of WR templates do not vary with environment. However, it is well known that SMC WN stars possess weak emission lines (Conti, Garmany & Massey 1989). In spite of small statistics and a large scatter, the mean He ii ��4686 line luminosity of WN2���4 stars in the LMC is 1035.9 ergs���1, a factor of five times higher than the mean of equivalent stars in the SMC (Crowther & Hadfield 2006). The signature of WN stars is most readily seen in star forming galaxies at He ii ��1640, where the dilution from other stellar types is at its weakest (e.g. Hadfield & Crowther 2006). The strongest UV, optical, and near-IR lines indicate flux ratios of I(He ii 1640)/I(He ii 4686)���10 and I(He ii 4686)/I(He ii 1.012��m)���6 for WN stars spanning SMC to Milky Way metallicities. Similar comparisons for WC stars are hindered because the only carbon-sequence WR stars at the low metallicity of the SMC and IC 1613 are WO stars. Their emission line fluxes are systematically weaker than WC stars in the LMC and Milky Way (Kingsburgh et al. 1995 Kingsburgh & Barlow 1995 Schaerer & Vacca 1998). The mean C iv ����5801-2 line luminosity of WC4 stars in the LMC is 1036.5 erg s���1 (Crowther & Hadfield 2006). Again, detection of WC stars is favoured via ultraviolet spectroscopy of C iv ��1550. For WC stars, the strongest UV, optical, and near-IR lines possess flux ratios of I(C iv 1548-51)/I(C iv 5801- 12)���6 and I(C iv 5801-12)/I(C iv 2.08��m)���15.

14 Paul A. Crowther 2.4 Binary statistics and masses The observed binary fraction amongst Milky Way WR stars is 40% (van der Hucht 2001), either from spectroscopic or indirect techniques. Within the low metallicity Magellanic Clouds, close binary evolution would be anticipated to play a greater role because of the diminished role of O star mass-loss in producing single WR stars. However, where detailed studies have been carried out (Bartzakos, Moffat & Niemela 2001 Foellmi, Moffat & Guerrero 2003ab), a similar binary fraction to the Milky Way has been obtained (recall Figure 3), so metallicity-independent LBV eruptions may play a dominant role. The most robust method of measuring stellar masses is from Kepler���s third law of motion, particularly for eclipsing double-lined (SB2) systems, from which the inclination may be derived. Orbital inclinations may also be derived from linear polarization studies (e.g. St-Louis et al. 1993) or atmospheric eclipses (Lamontagne et al. 1996). Masses for Galactic WR stars are included in the van der Hucht (2001) compilation, a subset of which are presented in Figure 4 together with some more recent results. WC masses span a narrow range of 9��� 16M���, whilst WN stars span a very wide range of ���10���83M���, and in some cases exceed their OB companion, i.e. q = MWR/MO 1 (e.g. WR22: Schweickhardt et al. 1999). WR20a (SMSP2) currently sets the record for the highest orbital- derived mass of any star, with ��� 83M��� for each WN6ha component (Rauw et al. 2005). As discussed above, such stars are H-rich, extreme O stars with strong winds rather than classical H-poor WN stars. They are a factor of two lower in mass than the apparent ��� 150M��� stellar mass limit (Figer 2005), such that still more extreme cases may await discovery. Spectroscopic measurement of masses via surface gravities using photospheric lines is not possible for WR stars due to

Wolf-Rayet Stars 15 their dense stellar winds. 2.5 Rotation velocities Rotation is very di���cult to measure in WR stars, since photospheric features ��� used to estimate v sin i in normal stars ��� are absent. Velocities of 200���500 km s���1 have been inferred for WR138 (Massey 1980) and WR3 (Massey & Conti 1981), although these are not believed to represent rotation velocities, since the former has a late-O binary companion, and the absorption lines of the latter are formed within the stellar wind (Marchenko et al. 2004). Fortunately, certain WR stars do harbour large scale structures, from which a rotation period may be inferred (St-Louis et al. 2007). Alternatively, if WR stars were rapid rotators, one would expect strong de- viations from spherical symmetry due to gravity darkening (Von Zeipel 1924 Owocki, Cranmer & Gayley 1996). Harries, Hillier & Howarth (1998) studied linear spectropolarimetric datasets for 29 Galactic WR stars, from which just four single WN stars plus one WC+O binary revealed a strong line effect, sug- gesting significant departures from spherical symmetry. They presented radiative transfer calculations which suggest that the observed continuum polarizations for these stars can be matched by models with equator to pole density ratios of 2���3. Of course, the majority of Milky Way WR stars do not show a strong linear polarization line effect (e.g. Kurosawa, Hillier & Schulte-Ladbeck (1999). 2.6 Stellar wind bubbles Ring nebulae are observed for a subset of WR stars. These are believed to represent material ejected during the RSG or LBV phases that is photo-ionized