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Wolf–Rayet star
Wolf–Rayet star
Wolf–Rayet star
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James Webb Space Telescope image of the Wolf–Rayet star WR 124 and the nebula M1–67 surrounding it. NIRCam and MIRI composite.

Wolf–Rayet stars, often abbreviated as WR stars, are a rare heterogeneous set of stars with unusual spectra showing prominent broad emission lines of ionised helium and highly ionised nitrogen or carbon. The spectra indicate very high surface enhancement of heavy elements, depletion of hydrogen, and strong stellar winds. The surface temperatures of known Wolf–Rayet stars range from 20,000 K to around 210,000 K, hotter than almost all other kinds of stars. They were previously called W-type stars referring to their spectral classification.

Classic (or population I) Wolf–Rayet stars are evolved, massive stars that have completely lost their outer hydrogen and are fusing helium or heavier elements in the core. A subset of the population I WR stars show hydrogen lines in their spectra and are known as WNh stars; they are young extremely massive stars still fusing hydrogen at the core, with helium and nitrogen exposed at the surface by strong mixing and radiation-driven mass loss. A separate group of stars with WR spectra are the central stars of planetary nebulae (CSPNe), post-asymptotic giant branch stars that were similar to the Sun while on the main sequence, but have now ceased fusion and shed their atmospheres to reveal a bare carbon-oxygen core.

All Wolf–Rayet stars are highly luminous objects due to their high temperatures—thousands of times the bolometric luminosity of the Sun (L☉) for the CSPNe, hundreds of thousands L☉ for the population I WR stars, to over a million L☉ for the WNh stars—although not exceptionally bright visually since most of their radiation output is in the ultraviolet.

The naked-eye star systems γ Velorum and θ Muscae both contain Wolf-Rayet stars, and two of the most massive known stars, BAT99-98 and R136a1 in 30 Doradus, are also Wolf–Rayet stars.

Observation history

[edit]
Crescent Nebula
WR 136, a WN6 star where the atmosphere shed during the red supergiant phase has been shocked by the hot, fast WR winds to form a visible bubble nebula

In 1867, using the 40 cm Foucault telescope at the Paris Observatory, astronomers Charles Wolf and Georges Rayet[1] discovered three stars in the constellation Cygnus (HD 191765, HD 192103 and HD 192641, now designated as WR 134, WR 135, and WR 137 respectively) that displayed broad emission bands on an otherwise continuous spectrum.[2] 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. E.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-integer quantum numbers were substituted. It was later shown that these lines resulted from the presence of helium, the chemical element having just been discovered in 1868.[3] 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.[4]

By 1929, the width of the emission bands was being attributed to Doppler broadening, and hence 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.[5] 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 nebulosity at all.[6]

In addition to helium, Carlyle Smith Beals identified emission lines of carbon, oxygen and nitrogen in the spectra of Wolf–Rayet stars.[7][8] 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.[9]

In 1969, several CSPNe with strong oxygen  VI (O VI) emissions lines were grouped under a new "O VI sequence", or just OVI type.[10] Similar stars not associated with planetary nebulae were described shortly after and the WO classification was adopted for them.[11][12] The OVI stars were subsequently classified as [WO] stars, consistent with the population I WR stars.[13]

The understanding that certain late, and sometimes not-so-late, WN stars with hydrogen lines in their spectra are at a different stage of evolution from hydrogen-free WR stars has led to the introduction of the term WNh to distinguish these stars generally from other WN stars. They were previously referred to as WNL stars, although there are late-type WN stars without hydrogen as well as WR stars with hydrogen as early as WN5.[14]

Classification

[edit]
WR 137 spectrum
Spectrum of WR 137, a WC7 star[15] and one of the three original WR stars (horizontal axis : wavelength in Å)

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. Initially simply referred to as class W or W-type stars,[16][17] the classification was then split into stars with dominant lines of ionised nitrogen (N III, N IV, and N V) and those with dominant lines of ionised carbon (C III and C IV) and sometimes oxygen (O III – O VI), referred to as WN and WC respectively.[18] The two classes WN and WC were further split into temperature sequences WN5–WN8 and WC6–WC8 based on the relative strengths of the 541.1 nm He II and 587.5 nm He I lines. Wolf–Rayet emission lines frequently have a broadened absorption wing (P Cygni profile) suggesting circumstellar material. A WO sequence has also been separated from the WC sequence for even hotter stars where emission of ionised oxygen dominates that of ionised carbon, although the actual proportions of those elements in the stars are likely to be comparable.[6] WC and WO spectra are formally distinguished based on the presence or absence of C III emission.[19] WC spectra also generally lack the O VI lines that are strong in WO spectra.[20]

The WN spectral sequence was expanded to include WN2–WN9, and the definitions refined based on the relative strengths of the N III lines at 463.4–464.1 nm and 531.4 nm, the N IV lines at 347.9–348.4 nm and 405.8 nm, and the N V lines at 460.3 nm, 461.9 nm, and 493.3–494.4 nm.[21] 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.[22]

The type WN1 was proposed for stars with neither N IV nor N V lines, to accommodate Brey 1 and Brey 66 which appeared to be intermediate between WN2 and WN2.5.[23] The relative line strengths and widths for each WN sub-class were later quantified, and the ratio between the 541.1 nm He II and 587.5 nm, He I lines was introduced as the primary indicator of the ionisation level and hence of the spectral sub-class. The need for WN1 disappeared and both Brey 1 and Brey 66 are now classified as WN3b. The somewhat obscure WN2.5 and WN4.5 classes were dropped.[24]

Classification of WN spectra
Spectral Type Original criteria[19] Updated criteria[24] Other features
WN2 N V weak or absent N V and N IV absent Strong He II, no He I
WN2.5 N  V present, N IV absent Obsolete class
WN3 N IV ≪ N V, N III weak or absent He II/He I > 10, He II/C IV > 5 Peculiar profiles, unpredictable N V strength
WN4 N IV ≈ N V, N III weak or absent 4 < He II/He I < 10, N V/N III > 2 C IV present
WN4.5 N IV > N V, N III weak or absent Obsolete class
WN5 N III ≈ N IV ≈ N V 1.25 < He II/He I < 8, 0.5 < N V/N III < 2 N IV or C IV > He I
WN6 N III ≈ N IV, N V weak 1.25 < He II/He I < 8, 0.2 < N V/N III < 0.5 C IV ≈ He I
WN7 N III > N IV 0.65 < He II/He I < 1.25 Weak P-Cyg profile He I, He II > N III, C IV > He I
WN8 N III ≫ N IV He II/He I < 0.65 Strong P-Cyg profile He I, He II ≈ N III, C IV weak
WN9 N III > N II, N IV absent N III > N II, N IV absent P-Cyg profile He I
WN10 N III ≈ N II N III ≈ N II H Balmer, P-Cyg profile He I
WN11 N III weak or absent, N II present N II ≈ He II, N III weak or absent, H Balmer, P-Cyg profile He I, Fe III present

The WC spectral sequence was expanded to include WC4–WC11, although some older papers have also used WC1–WC3. The primary emission lines used to distinguish the WC sub-types are C II 426.7 nm, C III at 569.6 nm, C III/IV 465.0 nm, C IV at 580.1–581.2 nm, and the O V (and O III) blend at 557.2–559.8 nm.[19] The sequence was extended to include WC10 and WC11, and the subclass criteria were quantified based primarily on the relative strengths of carbon lines to rely on ionisation factors even if there were abundance variations between carbon and oxygen.[20]

Classification of WC spectra
Spectral type Original criteria[19] Quantitative criteria[20] Other features
Primary Secondary
WC4 C IV strong, C II weak, O V moderate C IV/C III > 32 O V/C III > 2.5 O VI weak or absent
WC5 C III ≪ C IV, C III < O V 12.5 < C IV/C III < 32 0.4 < C III/O V < 3 O VI weak or absent
WC6 C III ≪ C IV, C III > O V 4 < C IV/C III < 12.5 1 < C III/O V < 5 O VI weak or absent
WC7 C III < C IV, C III ≫ O V 1.25 < C IV/C III < 4 C III/O V > 1.25 O VI weak or absent
WC8 C III > C IV, C II absent, O V weak or absent 0.5 < C IV/C III < 1.25 C IV/C II > 10 He II/He I > 1.25
WC9 C III > C IV, C II present, O V weak or absent 0.2 < C IV/C III < 0.5 0.6 < C IV/C II < 10 0.15 < He II/He I < 1.25
WC10 0.06 < C IV/C III < 0.15 0.03 < C IV/C II < 0.6 He II/He I < 0.15
WC11 C IV/C III < 0.06 C IV/C II < 0.03 He II absent

For WO-type stars the main lines used are C IV at 580.1 nm, O IV at 340.0 nm, O V (and O III) blend at 557.2–559.8 nm, O VI at 381.1–383.4 nm, O VII at 567.0 nm, and O VIII at 606.8 nm. The sequence was expanded to include WO5 and quantified based the relative strengths of the O VI/C IV and O VI/O V lines.[25] A later scheme, designed for consistency across classical WR stars and CSPNe, returned to the WO1 to WO4 sequence and adjusted the divisions.[20]

Classification of WO spectra
Spectral type Original criteria[19] Quantitative criteria[20] Other features
Primary Secondary
WO1 O VII ≥ O V, O VIII present O VI/O V > 12.5 O VI/C IV > 1.5 O VII ≥ O V
WO2 O VII < O V, C IV < O VI 4 < O VI/O V < 12.5 O VI/C IV > 1.5 O VII ≤ O V
WO3 O VII weak or absent, C IV ≈ O VI 1.8 < O VI/O V < 4 0.1 < O VI/C IV < 1.5 O VII ≪ O V
WO4 C IV ≫ O VI 0.5 < O VI/O V < 1.8 0.03 < O VI/C IV < 0.1 O VII ≪ O V

Detailed modern studies of Wolf–Rayet stars can identify additional spectral features, indicated by suffixes to the main spectral classification:[24]

  • h for hydrogen emission;
  • ha for hydrogen emission and absorption;
  • o for no hydrogen emission;
  • w for weak lines;
  • s for strong lines;
  • b for broad strong lines;
  • d for dust (occasionally vd, pd, or ed for variable, periodic, or episodic dust).[26]

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 "+OB" is used to indicate the presence of absorption lines in the spectrum likely to be associated with a more normal companion star, or "+abs" for absorption lines with an unknown origin.[24]

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-type star. There is a strong tendency for WNE stars to be hydrogen-poor while the spectra of WNL stars frequently include hydrogen lines.[19][27]

Spectral types for the central stars of planetary nebulae are qualified by surrounding them with square brackets (e.g. [WC4]).[19][28] They are almost all of the WC sequence with the known [WO] stars representing the hot extension of the carbon sequence. There are also a small number of [WN] and [WC/WN] types, only discovered quite recently.[29][30][31][32] Their formation mechanism is as yet unclear. Temperatures of the planetary nebula central stars tend to the extremes when compared to population I WR stars, so [WC2] and [WC3] are common and the sequence has been extended to [WC12]. The [WC11] and [WC12] types have distinctive spectra with narrow emission lines and no He II and C IV lines.[33][28]

Nova GK Persei
GK Persei (Nova Persei 1901), which showed Wolf–Rayet features in its spectrum[5]

Certain supernovae observed before their peak brightness show WR spectra.[34] This is due to the nature of the supernova at this point: a rapidly expanding helium-rich ejecta similar to an extreme Wolf–Rayet wind. The WR spectral features only last a matter of hours, the high ionisation features fading by maximum to leave only weak neutral hydrogen and helium emission, before being replaced with a traditional supernova spectrum. It has been proposed to label these spectral types with an "X", for example XWN5(h).[35] Similarly, classical novae develop spectra consisting of broad emission bands similar to a Wolf–Rayet star. This is caused by the same physical mechanism: rapid expansion of dense gases around an extremely hot central source.[6]

Slash stars

[edit]

The separation of Wolf–Rayet stars from spectral class O stars of a similar temperature depends on the existence of strong emission lines of ionised helium, nitrogen, carbon, and oxygen, but there are a number of stars with intermediate or confusing spectral features. For example, high-luminosity O stars can develop helium and nitrogen in their spectra with some emission lines, while some WR stars have hydrogen lines, weak emission, and even absorption components. These stars have been given spectral types such as O3If∗/WN6 and are referred to as slash stars.[36]

Class O supergiants can develop emission lines of helium and nitrogen, or emission components to some absorption lines. These are indicated by spectral peculiarity suffix codes specific to this type of star:

  • f for N iii and He ii emission
  • f* for N and He emission with N iv stronger than N iii
  • f+ for emission in Si iv in addition to N and He
  • parentheses indicating He ii absorption lines instead of emission, e.g. (f)
  • double parentheses indicating strong He ii absorption and N iii emission diluted, e.g. ((f+))

These codes may also be combined with more general spectral type qualifiers such as p or a. Common combinations include OIafpe and OIf*, and Ofpe. In the 1970s, it was recognised that there was a continuum of spectra from pure absorption class O to unambiguous WR types, and it was unclear whether some intermediate stars should be given a spectral type such as O8Iafpe or WN8-a. The slash notation was proposed to deal with these situations, and the star Sk−67°22 was assigned the spectral type O3If*/WN6-A.[37] The criteria for distinguishing OIf*, OIf*/WN, and WN stars have been refined for consistency. Slash star classifications are used when the Hβ line has a P Cygni profile; this is an absorption line in O supergiants and an emission line in WN stars. Criteria for the following slash star spectral types are given, using the nitrogen emission lines at 463.4–464.1 nm, 405.8 nm, and 460.3–462.0 nm, together with a standard star for each type:[36]

Classification of slash stars
Spectral type Standard star Criteria
O2If*/WN5 Melnick 35 N iv ≫ N iii, N v ≥ N iii
O2.5If*/WN6 WR 25 N iv > N iii, N v < N iii
O3.5If*/WN7 Melnick 51 N iv < N iii, N v ≪ N iii

Another set of slash star spectral types is in use for Ofpe/WN stars. These stars have O supergiant spectra plus nitrogen and helium emission, and P Cygni profiles. Alternatively they can be considered to be WN stars with unusually low ionisation levels and hydrogen.[38] The slash notation for these stars was controversial and an alternative was to extend the WR nitrogen sequence to WN10 and WN11[39] Other authors preferred to use the WNha notation, for example WN9ha for WR 108.[40] A recent recommendation is to use an O spectral type such as O8Iaf if the 447.1 nm He i line is in absorption and a WR class of WN9h or WN9ha if the line has a P Cygni profile.[36] However, the Ofpe/WN slash notation as well as WN10 and WN11 classifications continue to be widely used.[41]

A third group of stars with spectra containing features of both O class stars and WR stars has been identified. Nine stars in the Large Magellanic Cloud have spectra that contain both WN3 and O3V features, but do not appear to be binaries. Many of the WR stars in the Small Magellanic Cloud also have very early WN spectra plus high excitation absorption features. It has been suggested that these could be a missing link leading to classical WN stars or the result of tidal stripping by a low-mass companion.[42]

Nomenclature

[edit]
Carina Nebula around Wolf–Rayet star WR 22
WR 22 in the Carina Nebula

The first three Wolf–Rayet stars to be identified, coincidentally all with hot O-class companions, had already been numbered in the Henry Draper 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.[43] The first three catalogues were not specifically lists of Wolf–Rayet stars and they used only existing nomenclature.[44][45][46]

The fourth catalogue of Wolf-Rayet stars numbered them 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.[21] Neither of these numbering schemes remains in common use.

Modern WR catalogues

[edit]

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 galactic WR stars. These are again a numerical sequence from WR 1 to WR 158 in order of right ascension.[47]

Compiled in 2001, the seventh catalogue and its annex used the same numbering scheme and inserted 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.[19][48] Modern high volume identification surveys use their own numbering schemes for the large numbers of new discoveries.[49] A 2006 Annex was added to the seventh catalog.

In 2011, an online Galactic Wolf Rayet Catalogue was set up, hosted by the University of Sheffield. As of February 2025, it includes 709 stars.[50]

Other numbering schemes

[edit]

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"[51] 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.[52] As of 2018, 154 WR stars are catalogued in the LMC, mostly WN but including about twenty-three WCs as well as three of the extremely rare WO class.[42][53] 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.[54] 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[55][56][57]

In 2012, an IAU working group expanded the numbering system from the Catalogue of Galactic Wolf–Rayet stars so that additional discoveries are given the closest existing WR number plus a numeric suffix in order of discovery. This applies to all discoveries since the 2006 annex, although some of these have already been named under the previous nomenclature; thus WR 42e is now numbered WR 42-1.[58]

Properties

[edit]

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 600 Wolf–Rayets have been catalogued in our own Milky Way Galaxy.[19][48][49][50] 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.[59] It is expected that there are fewer than 1,000 WR stars in the rest of the Local Group galaxies, with around 166 known in the Magellanic Clouds,[42] 206 in the Triangulum Galaxy,[60] and 154 in the Andromeda Galaxy.[61]

Outside the local group, whole galaxy surveys have found thousands more WR stars and candidates. For example, in the M101 Group, over a thousand potential WR stars have been detected, from magnitude 21 to 25,[62] and astronomers hope to eventually catalog over ten thousand.[63] These stars are expected to be particularly common in the Wolf–Rayet galaxies named after them and in starburst galaxies.[64]

Their 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.[15] 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-type stars).[12]

Physical properties of galactic population I WN stars[27][65][66]
Spectral
type		 Temperature
(K)		 Radius
(R☉)		 Mass
(M☉)		 Luminosity
(L☉)		 Absolute
magnitude		 Example
WN2 141,000 0.89 16 280,000 −2.6 WR 2
WN3 85,000 2.3 19 220,000 −3.2 WR 46
WN4 70,000 2.3 15 200,000 −3.8 WR 1
WN5 60,000 3.7 15 160,000 −4.4 WR 149
WN5h 50,000 20 200 5,000,000 −8.0 R136a1
WN6 56,000 5.7 18 160,000 −5.1 CD Crucis
WN6h 45,000 25 74 3,300,000 −7.5 NGC 3603-A1
WN7 50,000 6.0 21 350,000 −5.7 WR 120
WN7h 45,000 23 52 2,000,000 −7.2 WR 22
WN8 45,000 6.6 11 160,000 −5.5 WR 123
WN8h 40,000 22 39 1,300,000 −7.2 WR 124
WN9h 35,000 23 33 940,000 −7.1 WR 102ea

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. The nitrogen seen in the spectrum of WNh stars is still the product of CNO cycle fusion in the core, but it appears at the surface of the most massive stars due to rotational and convectional mixing while still in the core hydrogen burning phase, rather than after the outer envelope is lost during core helium fusion.[14]

Physical properties of galactic population I WO/C stars[67]
Spectral
type		 Temperature
(K)[67]		 Radius
(R☉)[67]		 Mass
(M☉)[67]		 Luminosity
(L☉)[67]		 Absolute
magnitude		 Example
WO2 200,000 0.7 22 630,000 −1.7 WR 142
WC4 117,000 0.9 10 158,000 −3.28 WR 143
WC5 83,000 3.2 18 398,000 −4.87 Theta Muscae
WC6 78,000 3.6 18 501,000 −4.75 WR 45
WC7 71,000 4.0 17 398,000 −4.8 WR 86
WC8 60,000 6.3 18 398,000 −5.32 Gamma Velorum
WC9 44,000 8.7 13 251,000 −5.57 WR 104

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 place on those belonging to binary systems as a product of the collision of the stellar winds forming the pair,[19] as is the case of the famous binary WR 104; however this process occurs on single ones too.[15]

A few – roughly 10% – of the central stars of planetary nebulae, despite their much lower masses – typically ~0.6 M☉ – are 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.[68] These are now generally excluded from the class denoted as Wolf–Rayet stars, or referred to as Wolf–Rayet-type stars.[27]

Metallicity

[edit]

The numbers and properties of Wolf–Rayet stars vary with the chemical composition of their progenitor stars. A primary driver of this difference is the rate of mass loss at different levels of metallicity. Higher metallicity leads to high mass loss, which affects the evolution of massive stars and also the properties of Wolf–Rayet stars. Higher levels of mass loss cause stars to lose their outer layers before an iron core develops and collapses, so that the more massive red supergiants evolve back to hotter temperatures before exploding as a supernova, and the most massive stars never become red supergiants. In the Wolf–Rayet stage, higher mass loss leads to stronger depletion of the layers outside the convective core, lower hydrogen surface abundances and more rapid stripping of helium to produce a WC spectrum.

These trends can be observed in the various galaxies of the local group, where metallicity varies from near-solar levels in the Milky Way, somewhat lower in M31, lower still in the Large Magellanic Cloud, and much lower in the Small Magellanic Cloud. Strong metallicity variations are seen across individual galaxies, with M33 and the Milky Way showing higher metallicities closer to the centre, and M31 showing higher metallicity in the disk than in the halo. Thus the SMC is seen to have few WR stars compared to its stellar formation rate and no WC stars at all (one star has a WO spectral type), the Milky Way has roughly equal numbers of WN and WC stars and a large total number of WR stars, and the other main galaxies have somewhat fewer WR stars and more WN than WC types. LMC, and especially SMC, Wolf–Rayets have weaker emission and a tendency to higher atmospheric hydrogen fractions. SMC WR stars almost universally show some hydrogen and even absorption lines even at the earliest spectral types, due to weaker winds not entirely masking the photosphere.[69]

The maximum mass of a main-sequence star that can evolve through a red supergiant phase and back to a WNL star is calculated to be around 20 M☉ in the Milky Way, 32 M☉ in the LMC, and over 50 M☉ in the SMC. The more evolved WNE and WC stages are only reached by stars with an initial mass over 25 M☉ at near-solar metallicity, over 60 M☉ in the LMC. Normal single star evolution is not expected to produce any WNE or WC stars at SMC metallicity.[70]

Rotation

[edit]
Hubble Spies Vast Gas Disk around Unique Massive Star
Artist's illustration of gas disk around massive hydrogen-rich WR 122

Mass loss is influenced by a star's rotation rate, especially strongly at low metallicity. Fast rotation contributes to mixing of core fusion products through the rest of the star, enhancing surface abundances of heavy elements, and driving mass loss. Fast rotation causes stars to remain on the main sequence longer than slow-rotating stars, evolve more quickly away from the red supergiant phase, or even evolve directly from the main sequence to hotter temperatures for very high masses, high metallicity or very rapid rotation.

Stellar mass loss produces a loss of angular momentum and this quickly brakes the rotation of massive stars. Very massive stars at near-solar metallicity should be braked almost to a standstill while still on the main sequence, while at SMC metallicity they can continue to rotate rapidly even at the highest observed masses. Rapid rotation of massive stars may account for the unexpected properties and numbers of SMC WR stars, for example their relatively high temperatures and luminosities.[69]

Binaries

[edit]

Massive stars in binary systems can develop into Wolf–Rayet stars due to stripping by a companion rather than inherent mass loss due to a stellar wind. This process is relatively insensitive to the metallicity or rotation of the individual stars and is expected to produce a consistent set of WR stars across all the local group galaxies. As a result, the fraction of WR stars produced through the binary channel, and therefore the number of WR stars observed to be in binaries, should be higher in low metallicity environments. Calculations suggest that the binary fraction of WR stars observed in the SMC should be as high as 98%, although less than half are actually observed to have a massive companion. The binary fraction in the Milky Way is around 20%, in line with theoretical calculations.[71]

Nebulae

[edit]
LHA 115 - N76A
AB7 produces one of the highest excitation nebulae in the Magellanic Clouds.
Main article: Wolf–Rayet nebula

A significant proportion of WR stars are surrounded by nebulosity associated directly with the star, not just the normal background nebulosity associated with any massive star forming region, and not a planetary nebula formed by a post-AGB star. The nebulosity presents a variety of forms and classification has been difficult. Many were originally catalogued as planetary nebulae and sometimes only a careful multi-wavelength study can distinguish a planetary nebula around a low mass post-AGB star from a similarly shaped nebula around a more massive core helium-burning star.[70][72]

Wolf–Rayet galaxies

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Main article: Wolf–Rayet galaxy

A Wolf–Rayet galaxy is a type of starburst galaxy where a sufficient number of WR stars exist that their characteristic emission line spectra become visible in the overall spectrum of the galaxy.[73] Specifically a broad emission feature due to the 468.6 nm He ii and nearby spectral lines is the defining characteristic of a Wolf–Rayet galaxy. The relatively short lifetime of WR stars means that the starbursts in such galaxies must have occurred within the last few million years, and must have lasted less than a million years or else the WR emission would be swamped by large numbers of other luminous stars.[74]

Evolution

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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.

History

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Although Wolf–Rayet stars have been clearly identified as an unusual and distinctive class of stars since the 19th century,[75] the nature of these stars was uncertain until towards the end of the 20th century. Before the 1960s, even the classification of WR stars was highly uncertain, and their nature and evolution was essentially unknown. The very similar appearance of the central stars of planetary nebulae (CSPNe) and the much more luminous classical WR stars contributed to the uncertainty.[76]

By about 1960, the distinction between CSPNe and massive luminous classical WR stars was more clear. Studies showed that they were small dense stars surrounded by extensive circumstellar material, but not yet clear whether the material was expelled from the star or contracting onto it.[77][78] The unusual abundances of nitrogen, carbon, and oxygen, as well as the lack of hydrogen, were recognised, but the reasons remained obscure.[79] It was recognised that WR stars were very young and very rare, but it was still open to debate whether they were evolving towards or away from the main sequence.[80][81]

By the 1980s, WR stars were accepted as the descendants of massive OB stars, although their exact evolutionary state in relation to the main sequence and other evolved massive stars was still unknown.[82] Theories that the preponderance of WR stars in massive binaries and their lack of hydrogen could be due to gravitational stripping had been largely ignored or abandoned.[83] WR stars were being proposed as possible progenitors of supernovae, and particularly the newly-discovered type Ib supernovae, lacking hydrogen but apparently associated with young massive stars.[82]

By the start of the 21st century, WR stars were largely accepted as massive stars that had exhausted their core hydrogen, left the main sequence, and expelled most of their atmospheres, leaving behind a small hot core of helium and heavier fusion products.[84][85]

Current models

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Blue bubble in Carina
WR 31a is surrounded by a blue bubble created by a powerful stellar wind impacting material expelled during earlier stages of the star's life (ESA/Hubble & NASA Acknowledgement: Judy Schmidt)

Most WR stars, the classical population I type, are now understood as being a natural stage 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.[86] They have lost or burnt almost all of their hydrogen and are now fusing helium in their cores, or heavier elements for a very brief period at the end of their lives.[86]

Massive main-sequence stars create a very hot core which fuses hydrogen very rapidly via the CNO process and results in strong convection throughout the whole star. This causes mixing of helium 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 also show nitrogen enhancement at the surface at a very young age, caused by changes in the proportions of carbon and nitrogen due to the CNO cycle. The enhancement of heavy elements in the atmosphere, as well as increases in luminosity, create strong stellar winds which are the source of the emission line spectra. These stars develop an Of spectrum, Of* if they are sufficiently hot, which develops into a WNh spectrum as the stellar winds increase further. 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.

WR stars are likely to end their lives violently rather than fade away to a white dwarf. Thus every star with an initial mass more than about 9 times the Sun would inevitably result in a supernova explosion (with the exception of direct collapse[87]), many of them from the WR stage.[27][86][88]

A simple progression of WR stars from low to hot temperatures, resulting finally in WO-type stars, is not supported by observation. WO-type 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-type stars are only formed from the most massive main-sequence stars,[15] 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. It is still unclear whether the WO spectrum is purely the result of ionisation effects at very high temperature, reflects an actual chemical abundance difference, or if both effects occur to varying degrees.[86][89][90][91]

Schematic evolution of stars by initial mass (at solar metallicity) [citation needed]
Initial mass (M☉) Evolutionary sequence Supernova type
~250+ None[87]
~140– ~250 WNh-WNE-WO Ic/Pair-instability
120– ~140 WNh → WN → WC → WO Ic
60–120 O → Of → WNh ↔ LBV →[WNL] IIn [citation needed]
45–60 O → WNh → 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 B → RSG II-P

Key:

  • 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
  • WNh: WN plus hydrogen lines
  • WNL: "late" WN-class Wolf–Rayet star (about WN6 to WN11)
  • WNE: "early" WN-class Wolf–Rayet star (about WN2 to WN6)
  • WN/WC: Transitional (transitioning from WN to WC) Wolf–Rayet star (may be WN#/WCE or WC#/WN)
  • WC: WC-class Wolf–Rayet star
  • WO: WO-class Wolf–Rayet star

Wolf–Rayet stars form from massive stars, although the evolved population I stars have lost half or more of their initial masses by the time they show a WR appearance. 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.[92] High-mass stars are very rare, both because they form less often and because they have short lives. This means that Wolf–Rayet stars themselves are extremely rare because they only form from the most massive main-sequence stars and because they are a relatively short-lived phase in the lives of those stars. This also explains why type Ib/c supernovae are less common than type II, since they result from higher-mass stars.

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 show helium and nitrogen at the surface only a few thousand years after they form, possibly before they become visible through the surrounding gas cloud. 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.[93]

The difficulties of modelling the observed numbers and types of Wolf–Rayet stars through single star evolution have led to theories that they form through binary interactions which could accelerate loss of the outer layers of a star through mass exchange. WR 122 is a potential example that has a flat disk of gas encircling the star, almost 2 trillion miles wide, and may have a companion star that stripped its outer envelope.[94]

Supernovae

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Main article: supernova

It is widely suspected that many type Ib and type Ic supernova progenitors are WR stars, although no conclusive identification has been made of such a progenitor.

Type Ib supernovae lack hydrogen lines in their spectra. The more common type Ic supernovae lack both hydrogen and helium lines in their spectra. The expected progenitors for such supernova are massive stars that respectively lack hydrogen in their outer layers, or lack both hydrogen and helium. WR stars are just such objects. All WR stars lack hydrogen and in some WR stars, most notably the WO group, helium is also strongly depleted. WR stars are expected to experience core collapse when they have generated an iron core, and resulting supernova explosions would be of type Ib or Ic. In some cases it is possible that direct collapse of the core to a black hole would not produce a visible explosion.[95]

WR stars are very luminous due to their high temperatures but not visually bright, especially the hottest examples that are expected to make up most 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.[90] A possible progenitor star which has disappeared at the location of supernova iPTF13bvn may be a single WR star,[96] although other analyses favour a less massive binary system with a stripped star or helium giant.[97][98] The only other possible WR supernova progenitor is for SN 2017ein, and again it is uncertain whether the progenitor is a single massive WR star or binary system.[99]

In 2022 astronomers from the Gran Telescopio Canarias reported the first supernova explosion of a Wolf–Rayet star. SN 2019hgp was a type Icn supernova and is also the first in which the element neon has been detected.[100][101][102]

Examples

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Main article: List of Wolf-Rayet stars

By far the most visible example of a Wolf–Rayet star is γ2 Velorum (WR 11), which is a bright naked eye star for those located south of 40 degrees northern latitude, although most of the light comes from an O7.5 giant companion. 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 only other Wolf–Rayet star brighter than magnitude 6 is θ Muscae (WR 48), a triple star with two O class companions. Both are WC stars. The "ex" WR star WR 79a (HR 6272) is brighter than magnitude 6 but is now considered to be a peculiar O8 supergiant with strong emission. The next brightest at magnitude 6.4 is WR 22, a massive binary with a WN7h primary.[19]

The most massive and most luminous star currently known, R136a1, is also a Wolf–Rayet star of the WNh type that is still fusing hydrogen in its core. 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 WNh star such as VFTS 682 is found outside such clusters, probably having been ejected from a multiple system or by interaction with other stars.

An example of a triple star system containing a Wolf–Rayet binary is Apep. It releases huge amounts of carbon dust driven by their extreme stellar winds. As the two stars orbit one another, the dust gets wrapped into a glowing sooty tail.

All of the very hottest non-degenerate stars (the hottest few) are Wolf–Rayet stars, the hottest of which being WR 102, which seems to be as hot as 210,000 K, followed by WR 142 which is around 200,000 K in temperature. LMC195-1, located in the Large Magellanic Cloud, should have a similar temperature, but at the moment this temperature is unknown.

a giant smouldering star
HD 184738, also known as Campbell's star. This is actually a planetary nebula and the central star is an old low-mass star unlike the main class of massive Wolf–Rayet stars.[103]

HD 45166 has been described as the most magnetic massive star known and as the first magnetic known Wolf-Rayet star.[104]

Only a minority of planetary nebulae have WR type central stars, but a considerable number of well-known planetary nebulae do have them.

Planetary nebulae with WR type central stars[105]
Planetary nebula Central star type
NGC 2452 [WO1]
NGC 2867 [WO2]
NGC 5189 (Spiral Planetary Nebula) [WO1]
NGC 2371-2 [WO1]
NGC 5315 [WO4]
NGC 40 [WC8]
NGC 7026 [WO3]
NGC 1501 [WO4]
NGC 6751 [WO4]
NGC 6369 (Little Ghost Nebula) [WO3]
MyCn18 (Hourglass Nebula) [WC]–PG1159

See also

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References

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Further reading

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[edit]
Revisions and contributorsEdit on WikipediaRead on Wikipedia

Wolf–Rayet star

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Wolf–Rayet stars are evolved, massive stars in an advanced stage of , characterized by their intense stellar winds that drive extreme mass loss and produce distinctive broad emission-line spectra dominated by , , carbon, and oxygen. These winds strip away the stars' outer envelopes, exposing hot, helium-burning cores that make them among the hottest and most luminous objects in the galaxy, with surface temperatures typically ranging from 20,000 K to over 200,000 K and luminosities reaching up to a million times that of the Sun. Named after French astronomers Charles and Georges Rayet, who discovered the class in while observing peculiar spectra in Cygnus, these stars represent a brief, final phase lasting roughly 500,000 years before core collapse into supernovae. The spectral classification of Wolf–Rayet stars divides them into subtypes based on their emission lines: WN stars show prominent and lines, indicating nitrogen-rich atmospheres; WC stars exhibit carbon and oxygen features from carbon-burning stages; and rare WO stars display strong oxygen lines. These stars originate from progenitors with initial masses of at least 20–25 solar masses, evolving rapidly due to their high initial mass and subsequent mass loss rates that can exceed 10^{-5} solar masses per year—far surpassing those of main-sequence O stars. Their powerful winds create expanding bubbles of ionized gas and dust, observable as nebulae, and contribute to galactic chemical enrichment by ejecting processed elements into the . In , Wolf–Rayet stars are crucial for understanding massive star , binary interactions, and the endpoints of stellar life cycles, as they are primary candidates for progenitors of Type Ib and Ic core-collapse supernovae. Despite their rarity—only about 700 known in the (as of 2025)—they play a significant role in feedback processes that regulate in galaxies. Ongoing research, including observations from telescopes like Hubble and , continues to refine models of their winds, binarity, and potential as progenitors.

Discovery and History

Initial Observations

The discovery of Wolf–Rayet stars occurred in when French astronomers Charles and Georges Rayet, observing from the , identified three faint stars in the constellation Cygnus—later cataloged as HD 191765, HD 192103, and HD 192641—whose spectra displayed exceptionally broad and bright emission lines.[Wolf & Rayet (1867)] Using a 40-cm Foucault refractor equipped with a visual spectrograph, they noted these lines were far wider and more intense than the narrow absorption features typical of ordinary stellar spectra, suggesting unusual physical conditions such as high velocities or turbulent atmospheres.[Wolf & Rayet (1867)] These initial observations puzzled astronomers, as the emission lines did not match known gaseous nebulae or other stellar types observed at the time, leading to speculation about variable or explosive phenomena.[ & ] Subsequent analyses of these broad lines, spanning several angstroms, revealed expansion velocities exceeding 1,000 km/s, a feature unprecedented in early . In the early , American astronomer Edward C. Pickering at Observatory confirmed and expanded the class through systematic photographic , identifying additional Wolf–Rayet stars in constellations such as Aquila and Vela, including what became known as WR 1 and WR 2.[Pickering (1891)] By 1910, Pickering's surveys had cataloged over a dozen such objects, establishing their distinct across the sky and ruling out them being mere peculiarities of the Cygnus examples.[Pickering (1881)] Prior to 1920, estimates of distances and luminosities for these stars relied on limited data, including apparent magnitudes and early statistical methods, placing the original Cygnus trio at roughly 200–500 parsecs and implying luminosities thousands of times that of the Sun to account for their observed brightness despite faint visual appearance. These rough calculations, often tied to associations with nearby open clusters or nebulae, highlighted their extreme intrinsic brightness but were hampered by the absence of precise trigonometric parallaxes.

Key Developments

In the 1920s and 1940s, astronomers increasingly recognized Wolf–Rayet stars as evolved massive stars characterized by hydrogen-deficient atmospheres, a shift driven by detailed spectral analyses at institutions like Harvard Observatory. played a pivotal role through her contributions to and abundance studies, including the identification of dozens of Wolf–Rayet stars in the during early 20th-century spectral surveys and her analysis of their anomalous compositions indicative of envelope stripping. From the 1950s to the 1970s, theoretical advances in models emphasized mass loss as the primary mechanism transitioning massive stars into the Wolf–Rayet phase, with strong winds removing the hydrogen-rich outer layers to expose and heavier elements. Key works by researchers including Peter S. Conti established this framework, integrating observational data on wind velocities and integrating it with models to explain the observed spectral peculiarities. Virginia Trimble's reviews further synthesized these developments, highlighting how episodic or continuous mass loss shapes the late stages of massive star evolution. The 1980s and 2000s saw observational breakthroughs from the , which imaged circumstellar nebulae around Wolf–Rayet stars to reveal intricate structures formed by ejected material and interactions with ambient gas. For instance, Hubble observations of the (NGC 6888) around WR 136 provided constraints on ionization models and physical conditions within the nebula, confirming wind-blown origins. These images also uncovered evidence for binary companions in many systems, where colliding winds enhance mass loss and nebula morphology. The mission's Data Release 2 in the 2010s delivered precise parallaxes for over 380 Galactic Wolf–Rayet stars, yielding reliable distances and calibrations that refined population studies and evolutionary pathways. In recent years (2023–2025), the has contributed initial insights into dust formation around Wolf–Rayet stars through mid-infrared imaging. Observations of in 2023 revealed a distinctive halo of gas and dust with knotty, episodic structures, highlighting the star's role as an efficient dust producer despite its hydrogen deficiency. By 2025, JWST studies of the binary detected 17 expanding carbon-rich dust shells formed during periastron passages, demonstrating periodic dust creation via wind collisions and its implications for enrichment.

Classification

Spectral Features

Wolf–Rayet stars exhibit spectra dominated by broad, intense emission lines from highly ionized species, primarily , carbon, and , which distinguish them from other stellar classes. These lines, such as He II λ4686, C IV λ5801/08, and N V λ4604/20, originate in the extended, low-density stellar envelopes where recombination and processes dominate due to the stars' extreme conditions. The high effective temperatures of Wolf–Rayet stars, ranging from approximately 20,000 K to over 200,000 K, promote these high states, while the low atmospheric densities—resulting from rapid mass loss—contribute to the lines' exceptional widths, often spanning several thousand km/s. A hallmark of Wolf–Rayet spectra is the near absence or extreme weakness of Balmer lines, such as Hα and Hβ, indicating severe surface depletion through processes like mixing or ejection during advanced evolutionary stages. This deficiency contrasts sharply with the strong helium emission, particularly from He II, which serves as a primary diagnostic for the class. The emission lines are typically formed in optically thick , where the stellar ionizes the outflowing material, leading to a rich in forbidden and permitted transitions of heavy elements. Many prominent lines display P Cygni profiles, featuring a broad emission core superimposed with blue-shifted absorption troughs that reveal the of the stellar winds. These winds, driven by on ionized metals, achieve terminal velocities of 1,000 to 2,000 km/s or higher, with the absorption components extending to velocities reflecting the wind acceleration. Such profiles are evident in key ions like N V and C IV, providing direct evidence of the massive, dynamic envelopes enveloping these stars. Ultraviolet spectroscopy further elucidates the spectral traits of Wolf–Rayet stars, uncovering additional resonance lines such as those from O VI and Si IV that are less prominent in the optical range. These UV features, observed with instruments like the International Ultraviolet Explorer (IUE), enhance subtype differentiation by highlighting variations in and abundance patterns. For instance, the relative strengths of lines like He II λ1640 and C IV λ1550 help distinguish nitrogen-rich from carbon-rich subtypes in a single sentence.

Subtype System

The subtype system for Wolf–Rayet stars is a hierarchical scheme primarily based on the dominant broad emission lines arising from specific elements, reflecting surface abundances and conditions in their intense stellar winds. WN subtypes are defined by prominent (N III–V) and (He I–II) emission lines, with classifications spanning WN2 to WN11. The numerical scale from 2 to 11 denotes increasing levels from higher (lower numbers, hotter stars) to lower, where subtypes 2–5 (WNE, early) emphasize N V 4604–20 and He II 4686 , while 6–11 (WNL, late) highlight N III 4634–54 and He I 5876 . WC subtypes feature strong carbon (C III–V) and helium lines, with oxygen (O III–V) contributions in later types, classified as WC4 to WC9. The scale follows a similar pattern, with lower numbers indicating higher : early WC (WCE, 4–6) dominated by C IV 5801–12 , and late WC (WCL, 7–9) by C III 5696 and O III–V blends around 4600–50 . WO represents a rare subtype distinguished by highly ionized oxygen lines (O V–VI, notably O VI 3811–34 ) exceeding those in WC stars, with subtypes WO1 to WO4 assigned via the ratio of O VI to C IV line strengths, where WO1 signifies the highest . Surveys in the 2000s, leveraging , extended this system to dusty WC stars obscured at optical wavelengths by circumstellar material, employing line ratios of C IV 2.07–2.32 μm and He I 2.06 μm for subtype assignments aligned with optical criteria.

Nomenclature

Catalogues

The VIIth Catalogue of Galactic Wolf–Rayet Stars, compiled by K.A. van der Hucht in 2001, provides a comprehensive census of 227 Population I Wolf–Rayet (WR) stars in the , including improved equatorial coordinates, spectral subtypes, and BV photometry for known objects while incorporating 71 newly identified stars. This catalogue builds upon earlier compilations dating back to initial 19th-century discoveries, serving as a foundational reference for Galactic WR populations. An annex published in 2006 extended the catalogue by adding 71 more Galactic WR stars and candidates, primarily from infrared and spectroscopic follow-ups, bringing the total confirmed count to 298 objects, with detailed notes on binarity, variability, and associations with clusters or nebulae. Subsequent surveys have substantially expanded the known Galactic WR inventory. A 2025 study leveraging Data Release 3 (DR3) low-resolution spectra identified and spectroscopically confirmed 33 new WR stars (17 WC-type and 16 WN-type) among brighter candidates, enhancing the completeness of the census in obscured regions. Combined with other recent efforts, such as and optical monitoring, these additions have increased the total number of known Galactic WR stars to approximately 710 as of 2025. Extragalactic WR catalogues focus on nearby galaxies to study population differences. In the (M31), a 2023 narrow-band imaging and spectroscopic survey confirmed 19 new WR stars across targeted fields, predominantly fainter WN and WC subtypes, raising the known total in M31 to over 170. Astronomical databases facilitate access to these compilations. The service, maintained by the Centre de Données astronomiques de , hosts the full VIIth Catalogue (III/215) and its annex (J/A+A/458/453), along with integrated data from DR3 and other surveys, enabling cross-queries by coordinates, photometry, and spectral types.

Designation Schemes

Wolf–Rayet stars receive primary numerical identifiers in the form "WR" followed by a sequential number, a system introduced in early 20th-century catalogues and refined through successive compilations. This designation assigns numbers as a running sequence ordered primarily by increasing , reflecting the historical order of discovery and systematic surveys of the . The scheme originated with the first comprehensive lists in the 1940s and 1950s, such as those by H.D. Babcock and Merrill, and was formalized in the Sixth Catalogue (van der Hucht et al. 1981), with numbers up to WR 150 covering early-identified stars mostly in the solar neighborhood and . Subsequent updates, including the Seventh Catalogue, extended the sequence to over 220 confirmed galactic WR stars while preserving the original numbering for consistency across studies. For instance, WR 75 refers to a WN5-type star associated with the RCW 104 in Norma. In addition to WR numbers, many prominent Wolf–Rayet stars are catalogued under the Henry Draper (HD) system, which assigns unique identifiers based on equatorial coordinates and photometry for northern and southern stars brighter than about 9th magnitude, facilitating cross-referencing with broader stellar databases. This legacy from the early 20th-century Harvard Observatory survey remains widely used for bright examples, often paired with WR designations in spectroscopic analyses. A representative case is HD 147419, equivalently known as WR 75, a nitrogen-rich WR star illuminating its surrounding . Similarly, HD 93131 corresponds to WR 24, a highly luminous WN6ha star in the Trumpler 27 cluster. Southern hemisphere Wolf–Rayet stars, particularly those below declination -20°, frequently bear designations from the Bonner Durchmusterung (BD), an extensive 19th-century visual survey by Argelander and colleagues that mapped over 324,000 stars across the southern sky with positional and magnitude data. These BD numbers, formatted as BD +/−[declination zone]°[sequence number], predate modern photometry but provide enduring positional anchors for early-detected WR objects. For example, WR 140, a well-studied WC7 + O5 binary, is listed as BD +43°3571 in northern extensions of the survey, highlighting its role in colliding-wind research. Binary Wolf–Rayet systems that exhibit photometric or spectroscopic variability are often assigned variable star designations from the International Variable Star Index (VSX) or General Catalogue of Variable Stars (GCVS), using formats like "V" followed by a number within a constellation (e.g., V444 Cygni). These names arise when variability—due to eclipses, wind interactions, or pulsations—is confirmed, enabling coordinated observations across global networks. V444 Cygni, for instance, denotes the eclipsing WN5 + O6 binary also known as WR 139 and HD 193576, whose variations stem from its short 4.9-day period and colliding winds. Some of the brightest Wolf–Rayet stars receive informal or Flamsteed names based on their visibility to the , derived from historical constellation atlases like those of (1603) or Flamsteed (1725), which prioritize Greek letters for order. These proper names simplify reference in popular astronomy and early literature, often for standout examples in southern skies. Theta Muscae (θ Mus), a WC6 star and the second-brightest WR visible to the unaided eye at magnitude 5.5, exemplifies this, serving as WR 48 in systematic lists while anchoring studies of its associated emission nebulae.

Physical Properties

Composition and Spectra

Wolf–Rayet (WR) stars exhibit highly hydrogen-deficient envelopes, with the hydrogen-to-helium mass ratio typically H/He < 0.4, reflecting the stripping of outer layers and exposure of CNO-cycle processed material. This deficiency arises from intense mass loss, revealing atmospheres dominated by helium and heavy elements synthesized through hydrogen burning via the CNO cycle, where carbon, nitrogen, and oxygen act as catalysts, leading to enhanced abundances of nitrogen and helium while depleting hydrogen. Non-local thermodynamic equilibrium (non-LTE) atmospheric models confirm that these envelopes show surface compositions consistent with CNO equilibrium, with trace hydrogen present in some cases but overall low H content. In WN-type WR stars, the envelopes are primarily composed of helium and nitrogen, with nitrogen abundances elevated due to CNO processing, often reaching several times solar values while carbon remains depleted. These stars display strong emission lines from ionized helium (He II) and nitrogen ions (N III–V), indicative of their hot, nitrogen-rich winds. In contrast, WC-type WR stars feature envelopes rich in helium, carbon, and oxygen, with carbon mass fractions reaching up to 50–80% in advanced subtypes, resulting from further helium burning and triple-alpha processes that produce carbon. Oxygen lines are also prominent in WC spectra, particularly in the rarer WO subtype, where O abundances are higher relative to carbon. The winds of WR stars are clumped, with dense blobs embedded in a low-density interclump medium, leading to volume filling factors of approximately 0.1–0.3. This clumping influences the formation of broad emission line profiles by enhancing line strengths through higher local densities while reducing overall mass-loss rate estimates derived from smooth-wind models. Observations of line variability and polarization support this structure, with clumping factors implying density contrasts of 10–100, affecting the interpretation of spectral diagnostics. Electron temperatures in WR wind envelopes can be estimated from the observed intensity ratio of He II λ4686 to He I λ5876, where Te≈T_e \approx function of (observed He II λ4686 / He I λ5876 ratio), as this ratio is sensitive to ionization balance and temperature in the helium recombination zones. Such diagnostics, derived from non-LTE models, typically yield electron temperatures around 20,000–30,000 K in the line-forming regions.

Metallicity and Rotation

Wolf–Rayet (WR) stars at subsolar metallicities exhibit significantly weaker stellar winds compared to their solar-metallicity counterparts, primarily due to the reduced opacity from metal lines that drive radiatively accelerated outflows. This effect is prominently observed in WR populations within the Large Magellanic Cloud (LMC, Z ≈ 0.5 Z⊙) and Small Magellanic Cloud (SMC, Z ≈ 0.2 Z⊙), where WN-type WR stars display mass-loss rates and wind velocities lower by factors of 2–5 relative to Galactic analogs. At these low metallicities, the diminished wind strength hinders efficient envelope stripping, resulting in fewer WR stars overall and potentially altered evolutionary pathways, though single-star models suggest the WR phase duration may be extended in some rotating scenarios to compensate for reduced mass loss. Rotational velocities in WR stars are typically inferred from emission line profile distortions rather than direct photospheric measurements, given the dominance of extended winds. Equatorial rotation speeds generally range from 100 to 300 km s⁻¹, with projected values (v sin i) often around 200 km s⁻¹ for many objects, reflecting inheritance from progenitor O stars moderated by angular momentum loss during evolution. Higher velocities, up to ≈600 km s⁻¹, occur in select cases and manifest as asymmetries in spectral lines, such as skewed or broadened P Cygni profiles, due to rotational distortion of the wind geometry. Magnetic fields are detected in approximately 10% of WR stars through spectropolarimetric observations, with surface strengths on the order of a few kilogauss (kG) in confirmed cases, comparable to the incidence in massive O stars. These fields, often identified via Zeeman splitting in emission lines, are probed using instruments like FORS2 on the VLT, revealing organized structures in stars such as HD 45166 (B ≈ 43 kG, an outlier). Detection rates remain low due to the challenges of polarimetric signals in dense, ionized winds, but they indicate that magnetism may play a role in a minority of WR systems. The ratio of rotational velocity to terminal wind velocity (v_rot / v_∞, typically 0.05–0.3 given v_∞ ≈ 1000–3000 km s⁻¹) influences wind collimation in WR stars, particularly when combined with weak magnetic fields. At higher ratios, centrifugal forces and magnetic confinement can compress equatorial winds into disks or bipolar structures, leading to enhanced polarization and line profile variations that signal non-spherical outflows. This interplay is evident in hydrodynamical models, where v_rot / v_∞ > 0.1 promotes asymmetry without requiring binary interactions.

Mass Loss Rates

Wolf–Rayet stars exhibit exceptionally high mass loss rates, typically in the range of 10−510^{-5} to 10−410^{-4} M⊙M_\odot yr−1^{-1}, driven by their intense stellar winds that strip away the outer layers over short timescales. These winds reach terminal velocities between 1,000 and 3,000 km s−1^{-1}, significantly faster than those of main-sequence O stars, enabling the exposure of the hot . Such extreme ejection rates, observed across both nitrogen-rich (WN) and carbon-rich (WC) subtypes, play a crucial role in the stars' evolution and their chemical enrichment of the . The primary mechanism powering these winds is radiative driving, where ultraviolet photons from the star's hot surface (Teff≳30,000T_\mathrm{eff} \gtrsim 30,000 K) are absorbed and re-emitted by metal ions in the stellar atmosphere and wind, transferring momentum to the outflow. This process is particularly efficient in Wolf–Rayet stars due to their high luminosity and abundance of heavy elements, leading to optically thick winds with multiple scatterings. The overall momentum balance can be approximated by the relation M˙v∞=(L/c)Γ\dot{M} v_\infty = (L / c) \Gamma, where M˙\dot{M} is the mass loss rate, v∞v_\infty the terminal velocity, LL the stellar luminosity, cc the speed of light, and Γ\Gamma the force multiplier that quantifies the enhancement of radiative acceleration over Thomson scattering alone, often exceeding unity by factors of 10–100 depending on density and metallicity. However, theoretical models assuming homogeneous, smooth winds overestimate the true mass loss rates, as observations reveal significant clumping—small-scale enhancements—in Wolf–Rayet winds, arising from instabilities in the radiative acceleration. Clumping reduces the effective M˙\dot{M} by factors of 2–5 compared to smooth models, primarily because diagnostics like emission line strengths and radio fluxes scale with the square of , leading to underestimation if inhomogeneities are ignored. This adjustment is critical for accurate evolutionary modeling, as it implies less severe mass stripping than previously thought. Observationally, mass loss rates are derived from radio continuum observations, which probe free-free thermal emission from the ionized plasma, providing a direct measure largely insensitive to clumping if calibrated properly. Complementary estimates come from spectra, where the absorption troughs in P Cygni profiles of resonance lines (e.g., C IV or N ) yield both terminal velocities and, when combined with models, the M˙\dot{M} values. These methods have been refined through multi-wavelength campaigns, yielding consistent results for dozens of Galactic and Magellanic Cloud Wolf–Rayet stars.

Binary Systems

WR Binaries

Approximately 40% of known Galactic Wolf–Rayet (WR) stars are observed to reside in binary systems. These binaries typically consist of a WR star paired with a massive O- or early B-type companion, reflecting the evolutionary pathways of massive stars where the initially more massive component evolves into the WR phase. Orbital periods in these systems often range from a few days to several months, with WN-type WR binaries peaking at 1–10 days and WC types showing a broader distribution extending to longer periods. Detection of WR binaries relies on spectroscopic observations revealing radial velocity variations from the companion's absorption lines or the WR emission lines, as well as photometric monitoring for eclipses in close systems. Eclipsing binaries, such as V444 Cygni (WR 139, a WN5 + O6 system), provide particularly valuable insights, enabling direct measurements of stellar radii, orbital inclinations (typically near 78° for V444 Cygni), and component masses through modeling. These systems highlight the challenges in resolving the compact, hot WR cores amid strong stellar winds. In WR + O/B binaries, the powerful winds from both components collide, forming a region that generates hard emission through thermal processes in the post-shock gas. This colliding wind zone can also produce non-thermal radio emission and, in carbon-rich WC systems like , lead to periodic dust formation that manifests as rotating "pinwheel" structures observable in infrared imaging. Recent studies of newly discovered WR binaries, such as the 2024 analysis of a WN3 + O6 system, position these objects as archetypes for binary stripping mechanisms, where from the WR progenitor efficiently removes the envelope, shaping the system's toward potential binaries or double compact objects.

Slash Stars

Slash stars are Wolf–Rayet stars in close binary systems whose spectra display blended features from both the broad emission lines characteristic of the WR component and the absorption lines of an OB supergiant or main-sequence companion, often indicating unresolved orbital separations on the order of a few astronomical units. This composite appearance arises when the light from the two stars mixes without , leading to notations such as WN7+OB1 to describe the approximate spectral subtypes of the blended system. Approximately 20% of known Wolf–Rayet stars exhibit such composite spectra, underscoring the role of binarity in a significant portion of the WR population. Spectral dilution factors are used to quantify the relative flux contributions from the companion relative to the WR star, typically expressed as the ratio of the OB star's continuum flux to the total observed flux in specific bands, which weakens the apparent strength of WR emission lines. These factors, often derived from modeling the line profiles and continuum, enable the of the composite to estimate individual stellar parameters like temperatures and luminosities. For example, in the HD 211853 (also known as WR 153), classified as WN6+OB, the dilution factor reveals the OB companion's influence on the blended He II and N IV lines, allowing for refined spectral analysis. The blended spectra of slash stars carry evolutionary implications, as many such systems are products of where the WR progenitor, originally the more massive primary, has donated envelope material to the companion, accelerating its toward the WR phase through enhanced mass loss beyond stellar alone. This process can result in the WR star appearing as the secondary in the current system configuration. Unlike pure WR binaries, where high or photometric variability permits separation of the components' spectra, slash stars remain unresolved, complicating direct orbital parameter determination but providing indirect evidence of intimate binary interactions.

Associated Structures

Nebulae

Many Wolf–Rayet stars are enveloped by ring nebulae, which are ionized, expanding shells of gas formed through the interaction of the star's intense with circumstellar material ejected during prior evolutionary stages, such as the phase. These structures are primarily gas-dominated and trace the dynamical history of the star's loss. The high loss rates characteristic of Wolf–Rayet stars, often exceeding 10^{-5} M_\odot yr^{-1}, drive the sweeping up and compression of this ambient material into shell-like configurations. A prominent example is the (NGC 6888), associated with the WN6 –Rayet star WR 136 (HD 192163). This formed approximately 250,000 years ago when the fast wind from WR 136, reaching speeds over 1,500 km s^{-1}, collided with and energized the slower-moving from the star's preceding phase. NGC 6888 exhibits a bright, arc-shaped shell with intricate filaments, spanning about 7 pc in length and 5 pc in width, consistent with typical sizes of 1–5 pc for such Wolf–Rayet bubbles. The nebula expands at velocities of 55–110 km s^{-1}, aligning with the broader range of 20–50 km s^{-1} observed in many Wolf–Rayet ring nebulae. The intense ultraviolet radiation from the central Wolf–Rayet star ionizes the nebula, producing extended H II regions where hydrogen is singly ionized, alongside compact inner zones of He II emission where is doubly ionized due to the star's high exceeding 50,000 . In NGC 6888, spectroscopic observations reveal strong [S II] and [O III] lines indicative of shock-excited gas in the shell, with He II λ4686 emission tracing the hardest ionizing photons from WR 136. These ionization patterns highlight the nebula's role as a Strömgren sphere modified by wind compression, with electron densities around 10^3–10^4 cm^{-3}. Historical imaging from the has provided detailed views of these structures prior to the era, revealing clumpy, filamentary morphologies in the northeast rim of NGC 6888 that suggest ongoing interactions between the and dense gas clumps. These Wide Field Planetary Camera 2 observations, taken in 1999, resolved features as small as 0.05 pc, showing bright knots and bow shocks that illuminate the rapid of the over timescales of a few thousand years. Another illustrative case is the nebula M 1-67 surrounding the WN8 Wolf–Rayet star , a compact ring-like shell approximately 2 pc in diameter formed by similar wind interactions, expanding at around 150 km s^{-1} with a dynamical age of about 20,000 years. This structure exemplifies smaller, younger bubbles driven by runaway Wolf–Rayet stars, where the expansion velocity reflects the star's high peculiar motion of nearly 200 km s^{-1}.

Dust Shells

Carbon-rich Wolf–Rayet (WC) stars act as prolific dust factories, primarily through episodic mass ejections in binary systems where their dense winds collide with those of a companion star, creating shocked regions conducive to rapid . These collisions compress and cool the carbon-enriched material, enabling the formation of grains that radiate prominently in the infrared. of these systems confirms that the consists mainly of grains, with at least 19 WC stars exhibiting persistent dust production and 7 showing episodic activity tied to binary interactions. Recent James Webb Space Telescope (JWST) observations using the Mid-Infrared Instrument (MIRI) between 2023 and 2025 have unveiled intricate details of these dust shells around the WC8+O4 binary WR 140, revealing 17 concentric expanding shells visible in mid-infrared emission. These shells form periodically every ~7.9 years during periastron passages, when wind collisions intensify, and expand outward at velocities of 1–2% the speed of light (approximately 3,000–6,000 km s⁻¹), driven by the stellar winds and radiation pressure. The dust within these shells endures for centuries, contributing long-lived carbonaceous material to the interstellar medium and potentially influencing the formation of new stars and planets. Dust formation in these systems is inefficient, with only about 0.1%–1% of the stellar mass loss rate converting to dust, reflecting the brief windows of optimal conditions in the colliding winds. The resulting grains exhibit a range of temperatures from 300 K in cooler outer regions to 1,000 K near the binary, as derived from spectral energy distribution modeling of the infrared excess. Multi-shell structures are a hallmark of orbital modulation in these binaries, as evidenced by 2023 JWST MIRI imaging and dynamical modeling of WR 137, a WC7+O8 system, which resolved approximately 10 nested dust rings spanning over 130 years of episodic ejections. These shells arise from repeated dust formation events synchronized with the binary orbit, with expansion speeds around 1,700 km s⁻¹, highlighting how orbital geometry regulates dust output and spatial distribution.

Evolution

Formation Mechanisms

Wolf–Rayet (WR) stars primarily form through evolutionary pathways that expose their helium-burning cores by removing the overlying hydrogen-rich envelopes of massive stars. One key mechanism is the quasi-homogeneous , particularly prominent in low-metallicity environments, where rapid induces efficient mixing of nuclear-processed material from the core to the surface during the main-sequence phase. This process, driven by rotational velocities exceeding 200–300 km/s, leads to a gradual depletion of throughout the star without significant envelope stripping, resulting in WR-like spectral features (such as WNh subtypes) while the star is still hydrogen-burning. Observations of massive stars in the and support this pathway up to near-solar metallicities, as evidenced by the presence of hydrogen-rich WR stars with enhanced nitrogen and helium surface abundances. In higher-metallicity environments, single-star via enhanced mass loss becomes dominant for progenitors with initial masses exceeding approximately 25 M⊙. Strong radiative-driven winds during the post-main-sequence phases erode the , transitioning the into the WR phase as a -burning object with spectra dominated by broad emission lines of , , or . This mechanism is particularly effective for producing WN-type WR stars, where the stripping occurs primarily during the or stages, with mass-loss rates reaching 10^{-5} M⊙ yr^{-1} or higher. Seminal models indicate that progenitors around 40 M⊙ are required for WC subtypes, as further removal exposes carbon-oxygen burning products. Binary interactions provide an additional pathway, especially for lower-mass progenitors, through Roche-lobe overflow (RLOF) during mass transfer from the primary to the secondary star. In close binaries, the initially more massive donor star fills its Roche lobe during core hydrogen or helium burning, leading to unstable or stable mass transfer that strips the hydrogen envelope and exposes the helium core as a WR star. This process is crucial for explaining the observed fraction of WR stars in binary systems, which can constitute up to 40% of the population, and is supported by episodic mass-loss episodes during the RLOF phase that enhance wind strengths. Models of binary evolution highlight that this mechanism operates across a wide range of initial separations, producing WR companions with masses of 10–20 M⊙. Observational evidence for these formation mechanisms is bolstered by population synthesis studies that reproduce the observed ratios of WR stars to O-type stars (WR/O ratios) in galaxies like the Milky Way and Large Magellanic Cloud. These models, incorporating both single-star mass loss and binary channels, match empirical WR/O ratios of approximately 0.2–0.5 at solar metallicity when assuming initial binary fractions of 50–70% for massive stars. Such syntheses also align with luminosity functions, where the predicted distribution of WR subtypes peaks at log L/L⊙ ≈ 5.5–5.8, confirming the relative contributions of the pathways and their dependence on metallicity.

Models and Pathways

Computational models of Wolf-Rayet (WR) star evolution simulate the transition from the main-sequence phase to the WR phase by incorporating key physical processes such as rotation and mass loss. The evolutionary models, developed at the , provide a comprehensive framework for these simulations, computing tracks for stars with initial masses ranging from 20 to 120 M⊙M_\odot at various metallicities. These models account for rotational mixing, which enhances the transport of chemical elements to the surface, and enhanced mass-loss rates during the post-main-sequence phases, leading to the stripping of the hydrogen envelope and exposure of the core. In these tracks, massive stars enter the WR phase during core -burning, initially classified as WN types due to nitrogen-rich surfaces, and later transition to WC types as carbon is dredged up following helium exhaustion. The duration of the WR phase in these models typically spans 10510^5 to 10610^6 years, depending on the initial mass and , with higher-mass stars experiencing shorter phases due to intensified mass loss. An approximate expression for the WR lifetime is given by τWR≈MinitialM˙total,\tau_{\rm WR} \approx \frac{M_{\rm initial}}{\dot{M}_{\rm total}}, where MinitialM_{\rm initial} is the initial and M˙total\dot{M}_{\rm total} represents the cumulative mass-loss rate over the star's lifetime, highlighting the dominant role of mass loss in determining the phase length. Similar results emerge from other rotating models, such as those using the MESA code, which also cover initial masses of 25–120 M⊙M_\odot and confirm the WN-to-WC sequence through helium-burning evolution. Updates to these models in the have incorporated additional effects, such as surface fossil magnetic fields, which suppress mass loss and alter rotational evolution, potentially extending WR lifetimes and modifying surface compositions. Binary interactions, including and common-envelope evolution, are now integrated into extended frameworks like the Geneva binary models, influencing the pathways to the WR phase by accelerating envelope stripping in close systems. These advancements refine predictions for WR populations and their observational signatures.

Supernova Connections

Wolf–Rayet (WR) stars are widely regarded as the primary progenitors of Type Ib and Ic core-collapse supernovae (SNe Ib/c), owing to their advanced evolutionary stage where intense stellar winds have stripped away the hydrogen and, in some cases, helium envelopes, leaving exposed helium or carbon-oxygen cores. This hydrogen-deficient composition aligns with the spectral characteristics of SNe Ib/c, which lack hydrogen lines and, for Ic events, show minimal helium features. Evolutionary models indicate that WR stars reach core collapse after massive initial masses (typically 25–40 M⊙) and prolonged mass loss, transitioning from main-sequence O stars through luminous blue variable phases or binary interactions. A prominent example is SN 1998bw, a broad-lined Type Ic supernova discovered in 1998, whose progenitor properties—derived from pre-explosion imaging and light-curve analysis—match those of a WR star with an initial mass of approximately 25–40 M⊙. This event's high energy output (over 10^{52} erg) and association with gamma-ray burst GRB 980425 further highlight WR stars' role in energetic explosions, with the supernova's spectrum and circumstellar interaction signatures consistent with a WR wind environment. WR stars are also implicated in long-duration gamma-ray bursts (GRBs), particularly those collimated along the 's rotation axis, as in the case of GRB 980425, where the host galaxy exhibits Wolf–Rayet spectral signatures near the explosion site, supporting a massive WR . The "collapsar" model posits that rapidly rotating WR stars collapse to black holes, driving relativistic jets that produce GRBs while powering underlying SNe Ib/c. The extreme mass-loss rates of WR stars (∼10^{-5} M⊙ yr^{-1}) prior to explosion eject dense circumstellar material (CSM), which interacts with the supernova ejecta to produce bright radio, , and optical emission through shock heating and . This interaction explains the early light-curve peaks and spectral features in many SNe Ib/c, such as enhanced nitrogen lines from WR winds, and can amplify the explosion's by factors of 10–100. For the highest-mass WR stars (initial masses ≳60 M⊙), recent models predict that core collapse may result in fallback supernovae, where much of the falls back onto the proto-neutron star or , yielding dim or failed explosions with limited electromagnetic signatures. These outcomes suggest that such progenitors contribute to the mass function without bright SNe, consistent with the rarity of observed high-mass WR explosions.

Special Cases

Wolf–Rayet Galaxies

Wolf–Rayet galaxies are defined as a class of emission-line galaxies, often H II regions or starbursts, whose integrated optical spectra display prominent broad emission lines indicative of Wolf–Rayet (WR) stars, such as the He II λ4686 feature, signaling the presence of young, massive star clusters undergoing intense . These galaxies are characterized by recent bursts of , typically within the last 10 million years, where the collective spectra of WR populations dominate the integrated light. Such galaxies exhibit elevated star formation rates, commonly ranging from 10 to 100 M_⊙ yr⁻¹, driven by compact, dense clusters of massive stars, and they frequently occur in low-metallicity environments, including blue compact dwarf galaxies. A representative example is He 2-10, a nearby dwarf galaxy with a metallicity about 40% of solar, hosting a central starburst region with strong WR emission and a sustained high star formation rate. The low metallicity in these systems influences WR star evolution by reducing wind-driven mass loss, potentially extending their observable lifetimes compared to higher-metallicity counterparts. In Wolf–Rayet galaxies, the ratio of WR stars to late-type O stars (WR/L) is notably elevated compared to quiescent regions, often exceeding 0.1, due to the synchronized burst of that populates the WR phase more densely relative to O-star progenitors. This ratio, derived from equivalent widths of broad lines like He II λ4686, serves as a key diagnostic for burst intensity and is higher in systems with shorter, more vigorous episodes. Recent surveys in the 2020s, leveraging the (JWST), have enhanced detections of extragalactic WR populations in low-metallicity dwarf galaxies, revealing enrichment from WR winds in systems like the gravitationally lensed Sunburst Arc at z ≈ 2.37. These observations, using mid-infrared spectroscopy, confirm WR features in metal-poor environments and provide insights into early-universe analogs.

Notable Examples

One prominent example is , a WN8h-type Wolf–Rayet star known for its extreme with a of approximately 2,000 km/s, which drives the expansion of its surrounding M1-67. This , a bubble of ionized gas spanning about 6 light-years, results from the star's intense mass loss and is one of the youngest known Wolf–Rayet , with an estimated age of around 20,000 years. The star's high peculiar velocity of nearly 200 km/s relative to the local further shapes the into an asymmetric structure, highlighting WR 124's status as a runaway star. Theta Muscae, also designated WR 48, stands out as one of the brightest Wolf–Rayet stars visible from the , and the second-brightest overall, with an of about 5.5, making it a key target for southern observatories. Classified as a WC6 spectral type, it forms part of a triple system including O-class companions, contributing to its visual prominence despite the Wolf–Rayet component's intense emission lines. This system's luminosity and location in the constellation Musca have facilitated detailed studies of its wind interactions and surrounding emission nebulae. WR 104 exemplifies a binary Wolf–Rayet system with a WC9 primary and an O-type companion, renowned for its pinwheel-shaped formed by colliding stellar winds that periodically produce spirals. The system's orbital dynamics create a rotating plume observable in , raising concerns about a potential upon the stars' eventual supernovae, though recent observations indicate the beam is likely misaligned with by 30–40 degrees. This configuration provides a for understanding formation in massive binary . Recent observations in early 2025 have illuminated WR 140, a WC7 + O4–5 binary, revealing 17 concentric carbon-rich shells expanding at regular intervals due to periastron passages every eight years. These shells, imaged in mid-infrared, demonstrate episodic production from wind collisions, offering insights into interstellar enrichment processes. The system's high eccentricity and mass-loss rates underscore its role in modeling colliding-wind binaries.

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