A selfconsistent hydrodynamic calculation of a very massive star (MZAMS = 2OOM⊙) including turbulent pressure and energy has been performed. In the contraction phase after core hydrogen exhaustion, the star moves towards cool surface temperatures in the HR diagram (cf. Fig. 1). Consequently, (at Teff ⋍ 8000K) an envelope convection zone developes, and its inner boundery moves inwards with time. First, the envelope remains in hydrostatic equilibrium, with radiation pressure correspondingly decreasing as turbulent pressure increases (gas pressure is small). However, due to the fact, that the gradient of the turbulent pressure is directed inwards at the bottom of the convective zone, this part of the star rapidly contracts. Due to the released contraction energy, the luminosity locally exceeds the Eddington-luminosity. It cannot be transported outwards by convection in the upper part of the convection zone, where convective energy transport is inefficient (▽c ⋍ ▽r) . Thus, the local super-Eddington luminosity leads to the ejection of the overlying layers.

The implications of the nucleosynthetic activity of WR stars are reassessed, in view of recent experimental and observational data. It is confirmed that WR stars may 1) contribute significantly (up to ~20%) to the ~3 M⊙ of 26Al detected in the galactic plane through its 1.8 MeV line, 2) be responsible for the isotopic anomalies of 22Ne and 25,26Mg, detected in galactic cosmic rays (GCR), and 3) be responsible for the inferred presence of 26Al and 107Pd in the early solar system (and, perhaps, some other nuclei as well).

A theoretical interpretation of the observed upper luminosity limit is suggested here. Staritsin (1989) considered the core hydrogen and helium burning stages in a 64 M⊙ star. Mixing in a semi-convection zone in a diffusion approximation (Staritsin, 1987) and mass loss by stellar wind (de Jager et al., 1988) were taken into account. During MS evolution the star looses half of its initial envelope. After MS evolution an intermediate convective zone appears. The hydrogen content in the shell source increases. As a result, the star burns helium in its blue supergiant stage. After the hydrogen content in the envelope has decreased to 10% of its mass, the star looses mass with Wolf-Rayet mass loss rates according to de Jager et al. (1988). The star has a WR character during 5% of its full life time.

A direct comparison between the observed WR/WRprogenitor number ratio within 2.5 kpc from the sun and the predicted value (using evolutionary computations of single stars of Maeder and Meynet, 1987, A.&A.182, 243) reveals a discrepancy of at least a factor of two. In a previous study (Vanbeveren, 1990, A.&A. in press) I proposed a solution based on the incompleteness of the observed OB type star sample within 2.5 kpc from the sun. In this summary, I propose a theoretical explanation for the discrepancy. The theoretically predicted WR/WRprogenitor number ratio critically depends on the adopted M formalism in evolutionary computations during the red supergiant phase (RSG) of a massive star, especially in the mass range 20-40 M⊙. Since any M formalism predicts the mass loss rate with an uncertainty of at least a factor of two, I have tried to look for solutions for the WR/WRprogenitor problem by using different values of M during the RSG (in the mass range 20-40 M⊙); the M values and formalism that were adopted were always choosen within the observational uncertainty (i.e. within a factor of two when compared to the formalism used by Maeder and Meynet, 1987).

A search through literature reveals four methods in order to derive the mass of WR progenitors, i.e.

a. WR stars must be descendant from the most massive stars which share their galactic distribution,

b. the computation of detailed evolutionary models of massive close binaries up to the WR phase, able to explain the observational constraints of these WR binaries,

c. comparing the very narrow mass-luminosity relation of massive core helium burning stars predicted by evolution and estimated bolometric luminosities of WR members of stellar aggregates,

d. the minimum mass of the progenitor of a WR member of a cluster equals the mass of the most luminous star (or the star with the earliest spectral type) in the cluster.

Evolutionary computations of massive close binaries (MCB) including the effects of stellar wind (SW) and convective core overshooting predict that all massive primaries with ZAMS mass larger than 10 M⊙ start their core helium burning phase (CHeB) as bare helium cores; the hydrogen rich layers are removed on a timescale of the order of 104 yrs as a consequence of Roche lobe overflow (RLOF). The CHeB remnant after RLOF resembles closely a zero age CHeB star and its further evolution is entirely independent from its binary nature. Similarly as has been done previously by Vanbeveren and Packet (1979, A.&A.80, 242), I have performed a phenomenological study on the evolution of massive hydrogen less CHeB stars including the effect of SW mass loss using updated M determinations of van der Hucht et al. (1986, A.&A.168, 111). The SW mass loss rate formalism used in the computations is based on the following requirements:

The LBV phase is generally identified with the hydrogen shell burning phase of a star with initial mass larger than 40-50 M⊙ (see also Humphreys, this volume) and therefore in binaries with ZAMS component masses larger than 40-50 M⊙ and periods larger than 4-7 days the primary may experience a LBV mass loss phase before it reaches its critical equipotential surface (the Roche lobe). I then define the ‘LBV scenario’ of massive close binaries as follows:

when a binary component (initial mass larger than 40-50 M⊙) reaches the LBV phase prior to its Roche lobe overflow phase (RLOF), a stellar wind mass loss phase sets in at rates comparable to the rates encounterd during a RLOF process and which are large enough in order to prohibit the occurence of a RLOF.

Evolutionary models are computed for close binaries with initial primary masses larger than 50 M⊙ and mass ratios ranging between 0.2 and 1. Special attention is given at the predicted spectral type of the secondary component. In order to determine the mass of the primary at the end of its LBV phase, I have used the following general theorem holding for the most massive stars, i.e.

a hydrogen shell burning mass loser with mass larger than 50 M⊙ in a massive close binary restores thermal equilibrium (and becomes a WR star) when helium starts burning in its core and when its atmospherical hydrogen abundance has dropped to Xatm = 0.2-0.3 (by weight).

An energy distribution for P Cyg (B1 Ia+) has been produced using UKIRT photometry and spectra, flux-calibrated IUE high-resolution spectra, Johnson and Mitchell 13-colour photometry and IRAS photometry. Infrared excesses due to free-free emission from the wind have then been derived and modelled.

The distance of HR Carinae is determined with the reddening distance method, resulting into r ~ 5 kpc.

The LMC star R127 (=HDE 269858) has originally been classified as a late WN type (Walborn 1977). Since then, R127 has developed from spectral type B to A (Wolf et al. 1988) and is now categorized as a Luminous Blue Variable. R127 is embedded in gaseous material and a star cluster which are barely resolvable even on the highest-quality images available. Stahl (1987) studied the circumstellar material using narrow-band CCD images centered on nebular emission lines. No systematic census of the stars in the immediate vicinity of R127 has yet been published.

It is shown that a number of S Dor type stars in or near minimum brightness show small S Dor eruptions: δ VJ < 0m5

The main question in our understanding of the nature of W R stars is: which stars are their progenitors and how do they evolve to reach the WR stage? We have some ideas about progenitor's masses and, by mass luminosity relations, about their brightness when they are 0 stars. But, it is difficult to say anything for their possible Red phase.

The complete evolution of a star with an initial mass of 60 M⊙ and Z = Z⊙ (i.e. a typical W R progenitor) from the ZAMS through the supernova phase has been investigated. Mass loss in the different evolutionary stages, especially mass dependent W R mass loss, leads to a WO star (surface mass fractions {He,C,O} = {0.14, 0.38, 0.48}; cf. Fig. 1) of 4 . 2 M⊙ as pre-SN configuration. The low final mass may be typical for a wide range of initial masses (cf. Langer, 1989, Astr. Ap. 220, 135).

We argue that the progenitors of type Ib supernovae are moderately massive stars (M ~ 7 M⊙) in binary systems whereas the hypothesis that they originate from very massive stars (M > 20 M⊙) is not consistent with the observational evidence on SNIb.

We calculate the blue-red-blue evolution of the progenitor of SN 1987A in the HR diagram by adopting the Schwarzschild criterion for convection and by choosing appropriate parameters for mass loss and mixing [Fig. 1 (left) where metallicity is Z = 0.005]. During helium burning, the star moves from the blue to the red due to mass loss from 23 M⊙ to 16 M⊙. Figure 2 (right) shows the lifetime in effective temperature bins normalized to unity (solid) compared with the number of supergiants with −8 < Mboi < � 9 in the LMC, which is also normalized to unity (dashed). The model is qualitatively consistent with the observed histogram.

Surveys for Wolf-Rayet stars in nearby galaxies are briefly reviewed. The completeness and yield of these surveys are discussed in light of recent follow-up spectroscopy. A critical evaluation is made of our current knowledge of the Wolf-Rayet population in nearby galaxies, particularly the WC/WN ratio, the WR/O ratio, the WR surface density, and how these quantities vary within a galaxy and between galaxies, particularly as a function of metallicity. We compare the spectroscopic properties of Galactic and Magellanic Cloud WR stars with those in the more distant systems.