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 () for the CSPNe, for the population I WR stars, to 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
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, WR 135, and WR 137 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. 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. (The transition energies for an electron falling to the n=4 state around a nucleus of charge +2 are quite close to those of an electron falling to the n=2 state around a nucleus of charge +1, but with extra lines for electrons falling from states with odd values of n, whose energies correspond to transitions from a fictional n/2 state in hydrogen.) 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 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 nebulosity at all.
In addition to helium, Carlyle Smith Beals identified emission lines of carbon, oxygen and nitrogen 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.
In 1969, several CSPNe with strong oxygen VI (OVI) emissions lines were grouped under a new "OVI sequence", or just OVI type. Similar stars not associated with planetary nebulae were described shortly after and the WO classification was adopted for them. The OVI stars were subsequently classified as [WO] stars, consistent with the population I WR stars.
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.
Classification
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, the classification was then split 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 541.1 nm HeII and 587.5 nm HeI 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. WC and WO spectra are formally distinguished based on the presence or absence of CIII emission. WC spectra also generally lack the OVI lines that are strong in WO spectra.
The WN spectral sequence was 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.
The type WN1 was proposed for stars with neither NIV nor NV lines, to accommodate Brey 1 and Brey 66 which appeared to be intermediate between WN2 and WN2.5. The relative line strengths and widths for each WN sub-class were later quantified, and the ratio between the 541.1 nm HeII and 587.5 nm, HeI 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.
+Classification of WN spectra Spectral Type Original criteria Updated criteria Other features WN2 NV weak or absent NV and NIV absent Strong HeII, no HeI WN2.5 NV present, NIV absent Obsolete class WN3 NIV ≪ NV, NIII weak or absent HeII/HeI > 10, HeII/CIV > 5 Peculiar profiles, unpredictable NV strength WN4 NIV ≈ NV, NIII weak or absent 4 < HeII/HeI < 10, NV/NIII > 2 CIV present WN4.5 NIV > NV, NIII weak or absent Obsolete class WN5 NIII ≈ NIV ≈ NV 1.25 < HeII/HeI < 8, 0.5 < NV/NIII < 2 NIV or CIV > HeI WN6 NIII ≈ NIV, NV weak 1.25 < HeII/HeI < 8, 0.2 < NV/NIII < 0.5 CIV ≈ HeI WN7 NIII > NIV 0.65 < HeII/HeI < 1.25 Weak P-Cyg profile HeI, HeII > NIII, CIV > HeI WN8 NIII ≫ NIV HeII/HeI < 0.65 Strong P-Cyg profile HeI, HeII ≈ NIII, CIV weak WN9 NIII > NII, NIV absent NIII > NII, NIV absent P-Cyg profile HeI WN10 NIII ≈ NII NIII ≈ NII H Balmer, P-Cyg profile HeI WN11 NIII weak or absent, NII present NII ≈ HeII, NIII weak or absent, H Balmer, P-Cyg profile HeI, FeIII 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 CII 426.7 nm, CIII at 569.6 nm, CIII/IV 465.0 nm, CIV at 580.1–581.2 nm, and the OV (and OIII) blend at 557.2–559.8 nm. 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.
+Classification of WC spectra Spectral type Original criteria Quantitative criteria Other features Primary Secondary WC4 CIV strong, CII weak, OV moderate CIV/CIII > 32 OV/CIII > 2.5 OVI weak or absent WC5 CIII ≪ CIV, CIII < OV 12.5 < CIV/CIII < 32 0.4 < CIII/OV < 3 OVI weak or absent WC6 CIII ≪ CIV, CIII > OV 4 < CIV/CIII < 12.5 1 < CIII/OV < 5 OVI weak or absent WC7 CIII < CIV, CIII ≫ OV 1.25 < CIV/CIII < 4 CIII/OV > 1.25 OVI weak or absent WC8 CIII > CIV, CII absent, OV weak or absent 0.5 < CIV/CIII < 1.25 CIV/CII > 10 HeII/HeI > 1.25 WC9 CIII > CIV, CII present, OV weak or absent 0.2 < CIV/CIII < 0.5 0.6 < CIV/CII < 10 0.15 < HeII/HeI < 1.25 WC10 0.06 < CIV/CIII < 0.15 0.03 < CIV/CII < 0.6 HeII/HeI < 0.15 WC11 CIV/CIII < 0.06 CIV/CII < 0.03 HeII absent
For WO-type stars the main lines used are CIV at 580.1 nm, OIV at 340.0 nm, OV (and OIII) blend 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 sequence was expanded to include WO5 and quantified based the relative strengths of the OVI/CIV and OVI/OV lines. A later scheme, designed for consistency across classical WR stars and CSPNe, returned to the WO1 to WO4 sequence and adjusted the divisions.
+Classification of WO spectraSpectral type Original criteria Quantitative criteria Other features Primary Secondary WO1 OVII ≥ OV, OVIII present OVI/OV > 12.5 OVI/CIV > 1.5 OVII ≥ OV WO2 OVII < OV, CIV < OVI 4 < OVI/OV < 12.5 OVI/CIV > 1.5 OVII ≤ OV WO3 OVII weak or absent, CIV ≈ OVI 1.8 < OVI/OV < 4 0.1 < OVI/CIV < 1.5 OVII ≪ OV WO4 CIV ≫ OVI 0.5 < OVI/OV < 1.8 0.03 < OVI/CIV < 0.1 OVII ≪ OV
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; 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).
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.
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.