Wolf–Rayet stars (often referred to as WR stars) are evolved, massive stars (over 20 solar masses when they were on the main sequence) which are losing mass rapidly by means of a very strong stellar wind, with speeds up to 2000 km/s. They typically lose 10−5 solar masses a year, a billion times faster than the sun. Wolf–Rayet stars are extremely hot, with surface temperatures in the range of 30,000 K to around 200,000 K. They are also highly luminous, from tens of thousands to several million times the bolometric luminosity of the Sun, although not exceptionally bright visually since most of their output is in far ultraviolet and even soft X-rays.
In 1867, using the 40 cm Foucault telescope at the Paris Observatory, astronomers Charles Wolf and Georges Rayet discovered three stars in the constellation Cygnus (HD 191765, HD 192103 and HD 192641, now designated as WR 134, WR135, and WR137 respectively) that displayed broad emission bands on an otherwise continuous spectrum. Most stars only display absorption lines or bands in their spectra, as a result of overlying elements absorbing light energy at specific frequencies, so these were clearly unusual objects.
The nature of the emission bands in the spectra of a Wolf–Rayet star remained a mystery for several decades. Edward C. Pickering theorized that the lines were caused by an unusual state of hydrogen, and it was found that this "Pickering series" of lines followed a pattern similar to the Balmer series, when half-integral quantum numbers were substituted. It was later shown that the lines resulted from the presence of helium; a gas that was discovered in 1868. Pickering noted similarities between Wolf-Rayet spectra and nebular spectra, and this similarity led to the conclusion that some or all Wolf Rayet stars were the central stars of planetary nebulae.
By 1929, the width of the emission bands was being attributed to Doppler broadening, and hence that the gas surrounding these stars must be moving with velocities of 300–2400 km/s along the line of sight. The conclusion was that a Wolf–Rayet star is continually ejecting gas into space, producing an expanding envelope of nebulous gas. The force ejecting the gas at the high velocities observed is radiation pressure. It was well known that many stars with Wolf Rayet type spectra were the central stars of planetary nebulae, but also that many were not associated with an obvious planetary nebula or any visible nebulousity at all.
In addition to helium, emission lines of carbon, oxygen and nitrogen were identified in the spectra of Wolf–Rayet stars. In 1938, the International Astronomical Union classified the spectra of Wolf–Rayet stars into types WN and WC, depending on whether the spectrum was dominated by lines of nitrogen or carbon-oxygen respectively.
Wolf–Rayet stars were named on the basis of the strong broad emission lines in their spectra, identified with helium, nitrogen, carbon, silicon, and oxygen, but with hydrogen lines usually weak or absent. The first system of classification split these into stars with dominant lines of ionised nitrogen (NIII, NIV, and NV) and those with dominant lines of ionised carbon (CIII and CIV) and sometimes oxygen (OIII - OVI), referred to as WN and WC respectively. The two classes WN and WC were further split into temperature sequences WN5-WN8 and WC6-WC8 based on the relative strengths of the 541.1nm HeII and 587.5 nm HeI lines. Wolf-Rayet emission lines frequently have a broadened absorption wing (P Cygni profile) suggesting circumstellar material.
The WN spectral sequence has been expanded to include WN2 - WN9, and the definitions refined based on the relative strengths of the NIII lines at 463.4-464.1 nm and 531.4 nm, the NIV lines at 347.9-348.4 nm and 405.8 nm, and the NV lines at 460.3 nm, 461.9 nm, and 493.3-494.4 nm. These lines are well separated from areas of strong and variable He emission and the line strengths are well correlated with temperature. Stars with spectra intermediate between WN and Ofpe have been classified as WN10 and WN11 although this nomenclature is not universally accepted.
|Spectral Type||Criteria||Other emission lines|
|WN2||NV weak or absent||Strong HeII|
|WN2.5||NV present, NIV absent|
|WN3||NIV << NV, NIII weak or absent|
|WN4||NIV ≈ NV, NIII weak or absent|
|WN5||NIII ≈ NIV ≈ NV|
|WN6||NIII ≈ NIV, NV weak|
|WN7||NIII > NIV||Weak P-Cyg profile HeI, 468.6 nm HeII > NIII|
|WN8||NIII >> NIV||Strong P-Cyg profile HeI, 468.6 nm HeII ≈ NIII|
|WN9||NIII > NII, NIV absent||P-Cyg profile HeI|
|WN10||NIII ≈ NII||H Balmer, P-Cyg profile HeI|
|WN11||NIII weak or absent, NII present||H Balmer, P-Cyg profile HeI|
The WC spectral sequence has been expanded to include WC4 - WC9, although some older papers have also used WC1 - WC3. The WO types WO1 - WO4 have also been added for even hotter stars where emission of ionised Oxygen dominates that of ionised Carbon, although the actual chemical abundances in the stars are likely to be comparable. The primary emissions lines used to distinguish the WC sub-types are CII 426.7 nm, CIII at 569.6 nm, CIII/IV465.0 nm, CIV at 580.1-581.2 nm, and OV at 557.2-559.8 nm. For WO stars the main lines used are CIV at 580.1 nm, OIV at 340.0 nm, OV at 557.2-559.8 nm, OVI at 381.1-383.4 nm, OVII at 567.0 nm, and OVIII at 606.8 nm. The division between WC and WO spectra is easily made based on the presence or absence of CIII emission.
Detailed modern studies of Wolf Rayet stars can identify additional spectral features, indicated by suffixes to the main spectral classification:
- h for hydrogen emission;
- ha for hydrogen emission and absorption;
- w for wide lines;
- s for narrow (sharp) lines;
- d for dust (occasionally vd, pd, or ed for variable, periodic, or episodic dust).
The classification of Wolf Rayet spectra is complicated by the frequent association of the stars with dense nebulosity, dust clouds, or binary companions. A suffix of "+ abs" is often used to indicate the presence of absorption lines in the spectrum, likely to be associated with a more normal companion star.
The hotter WR spectral sub-classes are described as early and the cooler ones as late, consistent with other spectral types. WNE and WCE refer to early type spectra while WNL and WCL refer to late type spectra, with the dividing line approximately at sub-class six or seven. There is no such thing as a late WO star. There is a strong tendency for WNE stars to be hydrogen-poor while the spectra of WNL stars frequently include hydrogen lines.
The first three Wolf Rayet stars to be identified, coincidentally all with hot O companions, had already been numbered in the HD catalogue. These stars and others were referred to as Wolf–Rayet stars from their initial discovery but specific naming conventions for them would not be created until 1962 in the "fourth" catalogue of galactic Wolf Rayet stars. The first three catalogues were not specifically lists of Wolf Rayet stars and they used only existing nomenclature. The fourth catalogue numbered the Wolf Rayet stars sequentially in order of right ascension. The fifth catalogue used the same numbers prefixed with MR after the author of the fourth catalogue, plus an additional sequence of numbers prefixed with LS for new discoveries. Neither of these numbering schemes is in common use.
The sixth Catalogue of Galactic Wolf Rayet stars was the first to actually bear that name, as well as to describe the previous five catalogues by that name. It also introduced the WR numbers widely used ever since for all galactic WR stars. These are again a numerical sequence from WR 1 to WR 158 in order of right ascension. The seventh catalogue and its annex use the same numbering scheme and insert new stars into the sequence using lower case letter suffixes, for example WR 102ka for one of the numerous WR stars discovered in the galactic centre. Modern high volume identification surveys use their own numbering schemes for the large numbers of new discoveries.
Wolf Rayet stars in external galaxies are numbered using different schemes. In the Large Magellanic Cloud, the most widespread and complete nomenclature for WR stars is from the fourth Catalogue of Population I Wolf Rayet stars in the Large Magellanic Cloud, prefixed by BAT-99, for example BAT-99 105. Many of these stars are also referred to by their third catalogue number, for example Brey 77. A total of 134 WR stars is catalogued in the LMC, mostly WN but including three of the extremely rare WO class. Many of these stars are often referred to by their RMC (Radcliffe observatory Magellanic Cloud) numbers, frequently abbreviated to just R, for example R136a1.
In the Small Magellanic Cloud SMC WR numbers are used, usually referred to as AB numbers, for example AB7. There are only twelve known WR stars in the SMC, a very low number thought to be due to the low metallicity of that galaxy
Wolf–Rayet stars are a normal stage in the evolution of very massive stars, in which strong, broad emission lines of helium and nitrogen ("WN" sequence), carbon ("WC" sequence), and oxygen ("WO" sequence) are visible. Due to their strong emission lines they can be identified in nearby galaxies. About 500 Wolf–Rayets are catalogued in our own Milky Way Galaxy. This number has changed dramatically during the last few years as the result of photometric and spectroscopic surveys in the near-infrared dedicated to discovering this kind of object in the Galactic plane. It is expected that there are fewer than 1,000 WR stars in the rest of the Local Group galaxies, with around 150 known in the Magellanic Clouds, 206 in M33, and 154 in M31. Outside the local group, whole galaxy surveys have found thousands more WR stars and candidates, with particularly large numbers in starburst regions. For example, over a thousand WR stars have been detected in M101, from magnitude 21 to 25.
The characteristic emission lines are formed in the extended and dense high-velocity wind region enveloping the very hot stellar photosphere, which produces a flood of UV radiation that causes fluorescence in the line-forming wind region. This ejection process uncovers in succession, first the nitrogen-rich products of CNO cycle burning of hydrogen (WN stars), and later the carbon-rich layer due to He burning (WC and WO stars).
|Spectral Type||Temperature (K)||Radius||Mass||Luminosity||Absolute Magnitude|
It can be seen that the WNh stars are completely different objects from the WN stars without hydrogen. Despite the similar spectra, they are much more massive, much larger, and some of the most luminous stars known. They have been detected as early as WN5h in the Magellanic clouds.
|Spectral Type||Temperature (K)||Radius||Mass||Luminosity||Absolute Magnitude|
Some Wolf–Rayet stars of the carbon sequence ("WC"), especially those belonging to the latest types, are noticeable due to their production of dust. Usually this takes places on those belonging to binary systems as a product of the collision of the stellar winds forming the pair, as is the case of the famous binary WR 104; however this process occurs on single ones too.
A few (roughly 10%) of the central stars of planetary nebulae are, despite their much lower (typically ~0.6 solar) masses, also observationally of the WR-type; i.e., they show emission line spectra with broad lines from helium, carbon and oxygen. Denoted [WR], they are much older objects descended from evolved low-mass stars and are closely related to white dwarfs, rather than to the very young, very massive population I stars that comprise the bulk of the WR class. These are now generally excluded from the class denoted as Wolf–Rayet stars, or referred to as Wolf-Rayet type stars.
Theories about how WR stars form, develop, and die have been slow to form compared to the explanation of less extreme stellar evolution. They are rare, distant, and often obscured, and even into the 21st century many aspects of their lives are unclear.
Several astronomers, among them Rublev (1965) and Conti (1976) originally proposed that the WR stars as a class are descended from massive O-stars in which the strong stellar winds characteristic of extremely luminous stars have ejected the unprocessed outer H-rich layers. This has proved to be essentially correct, but with much complexity between a main sequence O star and the final WR star.
Early modelling of the evolution of massive stars showed that they evolve away from the main sequence, not towards hotter temperature and a WR state, but by expanding and cooling to become blue and then red supergiants. These supergiants are only modestly more luminous than the main sequence stars they originate from, but are progressively more unstable as their cores become hotter and their atmospheres more extended. Simple models of nuclear fusion showed that these red supergiants burned heavier elements in their cores until exploding as a supernova, but not becoming WR stars.
Further models showed that there was an upper limit to the stability of luminous stars. Sufficiently massive stars do not become red supergiants, instead shedding their atmospheres so quickly that they remain as blue supergiants, eventually shedding their atmospheres completely and entering the "Wolf Rayet funnel", an area of the HR diagram where WR stars become progressively smaller and hotter as they shed more and more of their outer layers. The suggestion was that earlier and hotter stars were the later stages of evolution from the later and cooler WR stars, but the results of this evolutionary sequence didn't match observations very well.
Most WR stars are now understood as being at a natural state in the evolution of the most massive stars (not counting the less common planetary nebula central stars), either after a period as a red supergiant, after a period as a blue supergiant, or directly from the most massive main sequence stars. Only the lower mass red supergiants are expected to explode as a supernova at that stage, while more massive red supergiants progress back to hotter temperatures as they expel their atmospheres. Some explode while at the yellow hypergiant or LBV stage, but many become Wolf Rayet stars.
Massive main sequence stars create a very hot core which fuses hydrogen via the CNO process and results in strong convection throughout the whole star. This causes mixing of fused elements to the surface, a process that is enhanced by rotation, possibly by differential rotation where the core is spun up to a faster rotation than the surface. Such stars show nitrogen at their surface at a very young age, combined with strong stellar winds. These stars develop an Of spectrum, Of* if they are sufficiently hot, which develops into a WNh spectrum as the levels of nitrogen at the surface increase. This explains the high mass and luminosity of the WNh stars, which are still burning hydrogen at the core and have lost little of their initial mass. These will eventually expand into blue supergiants (LBVs?) as hydrogen at the core becomes depleted, or if mixing is efficient enough (e.g. through rapid rotation) they may progress directly to WN stars without hydrogen.
Observations of supernova revealed that around a quarter of core collapse supernovae are of Type Ib, which originates from a progenitor with almost no hydrogen, and Type Ic, which originates from a progenitor with almost no hydrogen and very little helium. This corresponded rather well to WC and WO stars and as this was investigated it appeared that WR stars were likely to end their lives violently rather than fade away to a neutron star. Thus every star with an initial mass more than about 9 times the sun would inevitably result in a supernova explosion, many of them from the WR stage.
The simple progression of WR stars from low to hot temperatures, resulting finally in WO stars, is not supported by observation. WO stars are extremely rare and all the known examples are more luminous and more massive than the relatively common WC stars. Alternative theories suggest either that the WO stars are only formed from the most massive main sequence stars, and/or that they form an extremely short-lived end stage of just a few thousand years before exploding, with the WC phase corresponding to the core helium burning phase and the WO phase to nuclear burning stages beyond.
|Initial Mass (M☉)||Evolutionary Sequence||Supernova Type|
|60+||O → Of → WNLh ↔ LBV →[WNL]||IIn|
|45–60||O → WNLh → LBV/WNE? → WO||Ib/c|
|20–45||O → RSG → WNE → WC||Ib|
|15–20||O → RSG ↔ (YHG) ↔ BSG (blue loops)||II-L (or IIb)|
|8–15||O → RSG||II-P|
- O: O-type main-sequence star
- Of: evolved O-type showing N and He emission
- BSG: blue supergiant
- RSG: red supergiant
- YHG: yellow hypergiant
- LBV: luminous blue variable
- WNL: "late" WN-class Wolf–Rayet star (about WN6 to WN9)
- WNLh: WNL plus hydrogen lines
- WNE: "early" WN-class Wolf–Rayet star (about WN2 to WN6)
- WC: WC-class Wolf–Rayet star
- WO: WO-class Wolf–Rayet star
Although Wolf–Rayet stars form from exceptionally massive stars, most of them are only moderately massive because they only form after losing the bulk of their outer layers. For example, γ2 Velorum A currently has a mass around 9 times the sun, but began with a mass at least 40 times the sun. Higher-mass stars are much rarer, both because they form less often and because they only exist for a short time. This means that Wolf–Rayet stars themselves are very rare because they only form from the most massive main sequence stars, and explains why type Ibc supernovae are less common than type II. WNh stars, spectroscopically similar but actually a much less evolved star which has only just started to expel its atmosphere, are an exception and still retain much of their initial mass. The most massive stars currently known are all WNh stars rather than O-type main sequence stars, an expected situation because such stars start to move away from the main sequence only a few thousand years after they form. An alternative explanation is that these stars are so massive that they could not form as normal main sequence stars, instead being the result of mergers of less extreme stars.
Although it is widely accepted that most or all type Ibc supernovae progenitors were WR stars, no conclusive identification has been made of such a progenitor. WR stars are very luminous due to their high temperatures but not visually bright, especially the hottest examples that are expected to be supernova progenitors. Theory suggests that the progenitors of type Ibc supernovae observed to date would not be bright enough to be detected, although they place constraints on the properties of those progenitors. One candidate is under observation as pre-outburst observations show a likely WR star at the site of iPTF13bvn.
It is possible for a Wolf–Rayet star to progress to a "collapsar" stage in its death throes if it doesn't lose sufficient mass. This is when the core of the star collapses to form a black hole, either directly or by pulling in the surrounding ejected material. This is thought to be the precursor of a long gamma-ray burst. The compact object of Cygnus X-1 is one possible example.
The most visible example of a Wolf–Rayet star is Gamma 2 Velorum (γ² Vel), which is a naked eye star for those located south of 40 degrees northern latitude. Due to the exotic nature of its spectrum (bright emission lines in lieu of dark absorption lines) it is dubbed the "Spectral Gem of the Southern Skies". The second brightest is Theta Muscae. Both are multiple stars where the primary component is a Wolf Rayet type.
The most massive star and probably most luminous star currently known, R136a1, is also a Wolf–Rayet star of the WNh type indicating it has only just started to evolve away from the main sequence. This type of star, which includes many of the most luminous and most massive stars, is very young and usually found only in the centre of the densest star clusters. Occasionally a runaway Wolf–Rayet star such as VFTS 682 is found outside such clusters, probably having been ejected from a multiple system or by interaction with other stars.
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