A deep near-infrared spectroscopic survey of the Scutum-Crux arm for Wolf-Rayet stars††thanks: Based on observations with ESO telescopes at the La Silla Paranal Observatory under programme 094.D-0839(A)
We present an NTT/SOFI spectroscopic survey of infrared selected Wolf-Rayet candidates in the Scutum-Crux spiral arm (298, ). We obtained near-IR spectra of 127 candidates, revealing 17 Wolf-Rayet stars – a success rate – of which 16 are newly identified here. The majority of the new Wolf-Rayet stars are classified as narrow-lined WN5–7 stars, with 2 broad-lined WN4–6 stars and 3 WC6–8 stars. The new stars, with distances estimated from previous absolute magnitude calibrations, have no obvious association with the Scutum-Crux arm. Refined near-infrared (YHJK) classification criteria based on over a hundred Galactic and Magellanic Cloud WR stars, providing diagnostics for hydrogen in WN stars, plus the identification of WO stars and intermediate WN/C stars. Finally, we find that only a quarter of WR stars in the survey region are associated with star clusters and/or H ii regions, with similar statistics found for Luminous Blue Variables in the Milky Way. The relative isolation of evolved massive stars is discussed, together with the significance of the co-location of LBVs and WR stars in young star clusters.
keywords:stars: emission-line – stars: Wolf-Rayet – stars: evolution – galaxy: stellar content – infrared: stars.
Wolf-Rayet (WR) stars – the progeny of massive O-type stars – are excellent tracers of young stellar populations in galaxies owing to their unique spectroscopic signatures of strong, broad emission lines (Crowther, 2007). However, whilst WR surveys of nearby galaxies are nearing completeness (Massey et al., 2014), the Wolf-Rayet content of the Milky Way remains woefully incomplete (e.g. Shara et al., 2009) due to high dust obscuration at visual wavelengths. Our detailed knowledge of the evolution of massive stars remains unclear, with inaccuracies in earlier evolutionary phases magnified in the WR phase.
In addition, it is becoming clear that the conventional view of 20–25 O stars advancing through the Luminous Blue Variable (LBV) stage en route to the nitrogen- (WN) and carbon- (WC) sequence Wolf-Rayet phase and ultimately a stripped envelope core-collapse supernova (ccSN) is incomplete if not incorrect. First, a high fraction of massive stars are now known to be multiple (Sana et al., 2012), so the major effects of close binary evolution needs to be considered. Second, it has been proposed that LBVs are lower mass binary products, from an inspection of their spatial location in the Milky Way and Large Magellanic Cloud (LMC) with respect to Wolf-Rayet and O stars (Smith & Tombleson, 2015). Third, Sander et al. (2012) argue from a spectroscopic analysis of Milky Way WC stars that the most massive stars do not pass through this phase. Finally, it is not clear whether the most massive stars will undergo a bright SN explosion after core-collapse, since they may collapse directly to a black hole or produce a faint SN and fallback to a black hole (Langer, 2012).
Still, our Galaxy contains the largest spatially resolved population of WR stars, predicted to number (Rosslowe & Crowther, 2015a, b, hereafter RC15). The confirmed population has doubled over the previous decade, and currently stands at 111http://pacrowther.staff.shef.ac.uk/WRcat/. The Galactic disk therefore presents a rich hunting ground for further discoveries. Due to the large foreground interstellar dust extinction towards stars in the galactic disk, near and mid- infrared surveys are required.
The dense, ionized stellar wind of WR stars facilitate two approaches to infrared surveys. First, their strong, broad emission lines are amenable to near-IR narrow-band imaging (Shara et al., 2009; Shara et al., 2012; Kanarek et al., 2015). Second, their dense winds exhibit a free-free excess leading to unusual infrared colours, which have been exploited in the near-IR (Homeier et al., 2003; Hadfield et al., 2007) and mid-IR (Mauerhan et al., 2011; Messineo et al., 2012; Faherty et al., 2014). To date, the majority of spectroscopic follow-up has been carried out to an approximate depth of 11 mag. However, the (coarse) model of the Galactic Wolf-Rayet distribution developed by RC15 suggests follow-up spectroscopy is needed to 13 mag in order to sample the majority of Wolf-Rayet stars.
Here we exploit prior photometric approaches to spectroscopically survey a region of the Galactic disk to fainter limits than to date (13 mag). This has two interrelated goals: 1) the refinement and development of techniques that can be used to classify Wolf-Rayet stars using only infrared spectroscopy; 2) comprehensive searches for Wolf-Rayet stars in the Milky Way to allow more robust comparisons between their spatial locations and other massive stars; longer term goals involve the second data release (DR2) of Gaia which will provide parallaxes for hundreds of Wolf-Rayet stars, permitting their use as tracers of Galactic structure, and will be combined with upcoming large fibre-fed spectroscopic surveys, including WHT/WEAVE (Dalton et al., 2016) and VISTA/4MOST (de Jong et al., 2016).
This paper is structured as follows. In Section 2, we describe our photometric selection criteria and survey region, namely the Scutum-Crux spiral arm, the tangent to which lies at approximately (Georgelin & Georgelin, 1976). Spectroscopic observations of Wolf-Rayet candidates, plus some previously known Galactic WR templates, are presented in Section 3, including a brief description of non-WR stars. Refined near-IR classification criteria for Wolf-Rayet stars are presented in Section 4. Results for newly identified WR stars are presented in Section 5, including distance estimates. We consider the spatial location of Wolf-Rayet and other massive stars in Section 6, including discussion of prior inferences about the nature of Luminous Blue Variables. Finally, in Section 7 we reflect on the low success rate of the methodology employed, and share some motivating points for future IR surveys targeting WR stars.
2 Candidate Selection
Here we discuss our selection of sight lines towards the Scutum-Crux arm, plus our photometric criteria for the selection of candidate Wolf-Rayet stars. Specifically, we focus on , which Russeil et al. (2005) have previously highlighted in determining Galactic structure, since this intersects three proposed spiral arm features – Sagittarius-Carina, Scutum-Crux and Norma-Cygnus (Russeil, 2003, their Fig. 5). The majority of known Wolf-Rayet stars in this region lie at distances kpc, consistent with the nearby Sagittarius-Carina arm. However, assuming typical and for Galactic WR stars, it is possible to probe heliocentric distances kpc by identifying WR stars in the magnitude range mag. Therefore WR stars may provide a comparable and complimentary tracer of Galactic structure to commonly used non-stellar objects, i.e., H ii regions and atomic H gas. We confined our search to latitudes , ensuring that at the furthest expected distances, candidate WR stars remain within a few scale heights of the Galactic plane (FWHMpc for WRs, RC15).
|Number||– J2000 –||mag||mag||mag||mag||mag||mag||Grisms||Type|
We selected candidate Wolf-Rayet stars for which and from their near-IR (2MASS Skrutskie et al., 2006) and mid-IR (GLIMPSE-I Benjamin et al., 2003) photometry. We limited our survey to GLIMPSE-I point sources with a corresponing 2MASS detection222via the IPAC/NASA Infrared Science Archive: http://irsa.ipac.caltech.edu, requiring a minimum 2MASS quality flag of ‘C’ in the filter, and rejected sources with one or more Source Quality Flags in the GLIMPSE-I catalogue. We then used the TOPCAT333Available at: http://www.starlink.ac.uk/topcat/ tool to apply various cuts in colour and magnitude.
Mauerhan et al. (2011) identified several regions of colour space (their grey shaded region in Fig. 1) favoured by Wolf-Rayet stars, which we have adapted as follows:
In addition, Messineo et al. (2012) introduced additional reddening-free parameters and which we also utilise:
It is necessary to emphasise that not all WR stars occupy this parameter space, although there is no bias towards either WN or WC subtypes. Still, some dusty WC stars are offset from the majority of WR stars in the J-K vs colour colour diagram (Mauerhan et al., 2011, their Fig. 1), so our survey criteria are potentially biased against such stars. Approximately 250 sources satisfied these criteria. We subsequently cross-checked the co-ordinates of these with the SIMBAD444http://simbad.u-strasbg.fr/simbad/ database, to find any with previous identifications. Encouragingly, 14% of these were known WR stars (23% of known WRs in the survey area), 4% had non-WR classifications (mostly Be stars or young stellar objects), leaving 200 candidates with no previous spectral classification. Colour-colour diagrams for 191 candidates involving vs ) and vs are presented in Figure 1 together with reddening-free parameters Q1 and Q2 from Messineo et al. (2012).
|CTIO 4m/IRS||3000||Y||K||Mar 1996||a||C1|
|ESO 1.5m/B&C||Y||Feb 1982||b||E1|
|a: (Bohannan & Crowther, 1999), b: (Vreux et al., 1983); c: (Howarth & Schmutz, 1992), d: (Figer et al., 1997), e: (Hillier et al., 1983), f: (Crowther, 2000), g: (Homeier et al., 2003); h: (Crowther et al., 2006), i: This study; j: (Vreux et al., 1990), k: (Eenens et al., 1991), l: Crowther et al. (1995), m: (Crowther & Smith, 1996): n: (Tramper et al., 2015)|
3 NTT/SOFI spectroscopy of Wolf-Rayet candidates
We obtained near-IR spectroscopy of 127 WR candidates between 29–31 March 2015 (program ID 094.D-0839) using the Son-of-Isaac (SOFI) spectrograph at the New Technology Telescope (NTT). These represent 66% of the IR selected candidates presented in Fig. 1. Candidates were observed with the red grism (GR) covering the 1.53–2.52m spectral region, a dispersion of 10.2Å/pix, and a slit width of 1 arcsec, providing a spectral resolution of 600. All sources were observed using a standard ABBA sequence, including a small random offset in the A and B positions between exposures.
Before extracting 1D spectra, we subtracted a median dark frame from each individual frame, then subtracted adjacent AB pairs from one another. The result of this was 4 dark frame-corrected spectra for each source, free from sky lines. We extracted these 4 spectra for each object using IRAF. Wavelength calibration was performed using strong and isolated sky lines at known wavelengths (Rousselot et al., 2000) present in each raw frame, after which all spectra for each object were co-added.
Throughout each night, we periodically observed bright Vega-type telluric standard stars, at similar airmasses to the WR candidates. The removal of telluric spectral features was achieved using telluric in IRAF. We also used these telluric standards, together with Kurucz models of the same spectral types, to perform relative flux calibration, which are subsequently adjusted to match 2MASS photometry.
Of the 127 candidates, 17 stars were identified as WR stars. Of these, one candidate was subsequently matched to the recently discovered WN6 star 1093-1765 (= WR75-30) from Kanarek et al. (2015), such that 16 stars are newly identified as WR stars in this study. Previous surveys of Wolf-Rayet stars from IR photometric criteria have achieved similar efficiencies (Mauerhan et al., 2011; Faherty et al., 2014). We briefly discuss the nature of the non-WR stars in Sect. 3.1 and discuss the newly identified WR stars in Sect. 4. New WR stars are indicated in Fig. 1, with a subtype dependence apparent in the reddening-free Q1 vs Q2 diagram. Table 1 provides basic observational properties for the new Wolf-Rayet stars, for which we obtained NTT/SOFI spectroscopy, while Table B1 (available online) provides a list of all candidates for which we obtained spectroscopy, together with a brief note describing the nature of each source.
In addition to the candidate WR stars, we have also obtained NTT/SOFI spectroscopy of 14 Wolf-Rayet stars for which optical classifications have been undertaken, in order to refine near-IR based classification criteria (see Sect. 4). Finally, we also obtained blue grism (GB) spectroscopy with SOFI for the majority of newly identified WR stars. The GB observations cover the spectral region 0.95–1.64m, a dispersion of 7.0Å/pix, and an identical slit width of 1 arcsec, again providing 600. Data reduction was undertaken in an identical manner to the GR datasets.
3.1 Non Wolf-Rayet stars
A significant subset of the 110 candidates that were not confirmed to be WR stars exhibited a hydrogen emission line spectrum, with Br observed in 60 (55%) cases, plus often higher Brackett series (Br10, 11), and He i 2.058m emission present in a quarter of instances. These sources are likely to be massive young stellar objects (mYSOs) or Herbig AeBe stars (see e.g. Porter et al., 1998; Cooper et al., 2013). Br emission equivalent widths are typically 10–30Å, with He i/Br ratios of 0.3–1. The majority of Br emission lines are unresolved (FWHM30Å) although several stars (e.g. B#127, 147, 149) possess broad emission (FWHM50-60Å). Unusually, B#66 exhibits strong He i 2.058m emission, without significant Br emission, warranting follow up observations.
Mid-IR imaging has revealed circumstellar ring nebulae around many evolved stars (e.g., Wachter et al. 2010; Toalá et al. 2015). Such nebulae appear prominently in Spitzer 8.0µm images, owing to thermal emission from dust swept up by stellar winds. Mindful of this, we inspected Spitzer 8.0µm images centred on all candidates. One of the Br emission line sources, A#9 (SSTGLMC G330.7900-00.4539), is the central star of a striking oval mid-IR ring nebula, S65, which was identified by Churchwell et al. (2006) and studied by Simpson et al. (2012).
Strong absorption lines in the Brackett series are observed in one candidate, B#3, indicating an A- or late-B type star, with Br absorption observed in another object, B#54, albeit without other prominent features suggesting an early type star in this instance. Four candidates – B#21, B#100, B#122 and B#153 – exhibit prominent CO 2.3m bandhead absorption features, although none of these involve Br emission line sources, indicating a late-type star origin. The remaining 44 candidates (40% of the non WR stars) either have no prominent absorption or emission features, or the S/N achieved was insufficient to identify their nature.
4 Near-IR classification of Wolf-Rayet stars
The switch from optical to near-IR spectroscopy for the overwhelming majority of new Galactic Wolf-Rayet stars requires a reassessment of spectral classification criteria. Vreux et al. (1990) provided an classification scheme based upon Y-band observations of northern WR stars, while Eenens et al. (1991) devised a near-IR scheme for WC stars from 1–5m spectroscopy. More recently Crowther et al. (2006, hereafter C06) provide near-IR classification diagnostics for WN and WC stars, based on equivalent width ratios. Qualitatively, early-type WC4–6 stars possess broader emission lines than later WC7–9 subtypes, although exceptions do exist (Eenens & Williams, 1994).
An updated a quantitative near-IR classification of Wolf-Rayet stars is made feasible by access to a greatly expanded sample of Galactic and Magellanic Cloud WR stars for which optical classifications have been made, primarily Smith et al. (1996) for WN stars and Smith et al. (1990) for WC stars. The datasets utilised were drawn from various sources, primarily NTT/SOFI and UKIRT/CGS4, as summarised in Table 2. We have also inspected high resolution, intermediate resolution IRTF/SpeX spectroscopy of northern Galactic WR stars, provided by W.D. Vacca (priv. comm.), although we focus our criteria on moderate resolution (R = 600 – 1000), modest signal-to-noise spectroscopy in the 0.95 – 2.5m near infrared.
To date, no criteria for the identification of WN/C or WO stars from near-IR spectroscopy have been considered, nor has an attempt to distinguish between H-rich and H-deficient WN stars, although C06 did separate broad-lined WN stars (FWHM He ii 1.01µm 65Å) from narrow-lined counterparts. In our revised near-IR classification scheme we attempt to utilise pairs of lines from adjacent ionization stages of helium for WN stars, and adjacent ionization stages of carbon for WC stars. In some instances nitrogen lines are required for WN stars, in common with C06, plus ratios of carbon to helium lines are utilised for WN/C and WC stars, which will also depend upon their relative abundances. We omit from our discussion the near-IR classification of transition Of/WN stars, which has been considered by Bohannan & Crowther (1999) and Crowther & Walborn (2011). WN stars with intrinsic absorption features (WNh+abs, WNha) also offer specific challenges which will need to be considered separately.
A detailed description of the updated classification scheme is provided in Appendix A (available online), while we present a summary of Y, J, H and K-band classification diagnostics for WN, WN/C, WC and WO subtypes in Table 3. In Figs A1–A2 (available online) we present YJ-band and HK-band spectroscopy of template (optically classified) WN, WN/C, WC and WO stars, with line measurements provided in Tables A1–A3, also online. Overall, the use of solely IR diagnostics provide satisfactory classifications, although confidence in resulting spectral types requires multi-band spectral coverage, a minimum spectral resolution of 500 and moderate signal-to-noise. In many instances, observations at a single band prevent a refined classification. For example, WC4–7 stars may not be distinguished using solely K-band spectroscopy, while it is not possible to differentiate between broad-lined WN4–7 stars on the basis of low S/N spectroscopy in the Y-band. Still, reliable, Wolf-Rayet subtypes can be obtained from complete 1–2.5m spectroscopy, with the exception of broad-lined WN4–6 stars and WC5–6 stars. In addition, WO stars have a very distinctive near-IR spectrum, and WN/C stars possess characteristics in each of Y, J, H and K-bands which distinguish them from normal WN stars. In addition, the presence of hydrogen in WN stars can be identified in most subtypes, although very late subtypes are challenging since a low He ii 1.163/P or He ii 2.189/Br ratio may indicate either a high hydrogen content or a low ionization atmosphere.
4.1 Robustness of near-IR classification
In order to assess the robustness of the new scheme, we reclassify several WN and WC stars which have been discovered and classified from red optical spectroscopy. We utilise NTT/SOFI spectroscopy of four WN stars, WR62a, WR68a, WR93a from Homeier et al. (2003) (dataset N2 in Table 2), plus WR75-30 from our own observations (dataset N5 in Table 2), together with three WC stars, WR107a from Homeier et al. (2003) plus WR75aa and WR75c from our own observations. Near-IR spectra are presented in Figs 2–3. Individual line measurements are provided in Table 4 and 5 for WN and WC stars, respectively, while line ratios are presented in Table 6. Measurements have employed Gaussian fits, using the elf suite of commands within the Starlink DIPSO package555Available at: http://starlink.eao.hawaii.edu/starlink.
4.1.1 WN stars
WR62a was classified as WN5o by Shara et al. (1999, their source #11), and we support its classification as a narrow-lined WN star. Consequently, the primary diagnostics are the He i 1.08/He ii 1.01m ratio and K-band morphology. The former indicates a WN6 subtype, while He i + N iii 2.11 N v 2.10, favours a WN5–6 subtype. The P/He ii 1.16m ratio indicates WR62a is hydrogen-free, while the Br/He ii 2.19m suggests a borderline o/(h) classification, so overall we favour WN6o for WR62a. The same arguments and ratios apply to WR68a for which Shara et al. (1999, their source #13) assigned WN6o. We support this classification owing to its morphological similarity to WR62a.
WR93a (Th 3–28), was originally classified as WN2.5–3 by Acker & Stenholm (1990) and revised to WN6 by Miszalski et al. (2013) from optical spectroscopy. This is also a narrow-lined WN star so again we focus on its He i 1.08/He ii 1.01m ratio and K-band morphology. Both favour a WN6 subtype, with a significant hydrogen content from our multiple diagnostics (darker shaded regions in Fig. A4, available online), so we adopt WN6h for WR93a.
Kanarek et al. (2015, their source 1083-1765) originally classified WR75-30 as a WN6 star from near-IR spectroscopy. The He i 1.70/He ii 1.69m ratio favours a WN7 subtype, as does (He i + N iii 2.11)/He i 2.19m 1, while the Br/He ii 2.19m ratio lies in the hydrogen-free region of Fig. A4 so we favour WN7o for this star.
4.1.2 WC stars
WR75aa and WR75c were identified as WC9 stars by Hopewell et al. (2005) from red optical spectroscopy. All our primary near-IR diagnostics support this assessment, as do the secondary criteria involving helium for WR75c. WR75aa has a borderline WC8–9 classification from the He i-ii 1.7/C iv 1.74m ratio, but overall both stars are unambiguous WC9 stars.
Finally, WR107a (#18 from Shara et al., 1999) was originally classified as a WC6 star from red optical spectroscopy. Our primary criteria indicate the follow for WR107a: WC61 from both C iii 1.20/C iv 1.19 and C iii 2.11/C iv 2.07, WC5–8 from C ii 0.99/C iii 0.97. Our secondary criteria indicate WC5 from He i 1.08/He ii 1.01, and WC7 from C iii 0.97/He ii 1.01m (H-band spectra are unavailable), so although WC6 is plausible we provide a more cautious WC5–7 classification. Indeed, the primary optical diagnostic ratio (C iii 5696/C iv 5808) also favoured WC6–7 according to Shara et al. (1999).
In general, the K-band is preferred to shorter wavelengths for classification of highly reddened WR stars, but K-band spectral features of dusty WC stars are often masked by host dust emission. Extremely high S/N is required to identify K-band spectral features of the Quintuplet Q stars. By way of example, Liermann et al. (2009) assign a WC8/9d+OB subtype to Q3 (WR102ha) from K-band spectroscopy, whereas WC features are relatively prominent in deep H- and J-band spectroscopy. We confirm a WC9d subtype for Q3 on the basis of high S/N Gemini spectroscopy presented by Najarro et al. (2015), owing to C iii 1.20/C iv 1.191, C ii 1.78/C iv 1.741 and C iii 2.11/C iv 2.071.
|Narrow-lined WN stars (FWHM(He ii 1.01m) 1900 km s and (He ii 1.01m/Å) 2.5)|
|WN||He ii||log (He i 1.08/||log (P/||(He i 1.70/||(Br/||Notes||Templates|
|Subtype||1.012m||He ii 1.01)||He ii 1.16)||He ii 1.69)||He ii 2.19)|
|9||300||1.4||1.5 (h)||1.4||1.5 (h)||He i+N iii 2.11 He ii 2.19||WR105, BAT76|
|8||700||10.4||0.3 (o); 0.1 (h)||0.90.5||0.5 (o); 0.4 (h)||He i+N iii 2.11 He ii 2.19||WR40, WR123|
|7||800||0.40.2||0 (o); 0 (h)||0.20.2||0.1 (o); 0.1 (h)||He i+N iii 2.11 He ii 2.19||WR78, WR120|
|6||1200||0.00.2||–0.1 (o); –0.1 (h)||–0.1||–0.1 (o); –0.1 (h)||He ii 2.19 He i+N iii 2.11 N v 2.10||WR115|
|5||1400||–0.30.1||–0.2 (o); –0.2 (h)||–0.3||–0.2 (o); –0.2 (h)||He ii 2.19 He i+N iii 2.11 N v 2.10||BAT122|
|4||1500||–0.70.3||–0.3 (o); –0.3 (h)||–0.6||–0.3 (o); –0.3 (h)||He ii 2.19 N v 2.10 He i+N iii 2.11)||WR128, BAT75|
|3||1600||–1.0||–0.4 (o); –0.4 (h)||–0.8||–0.4 (o); –0.4 (h)||He ii 2.19 N v 2.10 He i+N iii 2.11||WR46, WR152|
|Broad-lined WNb stars (FWHM(He ii 1.01m) 1900 km s and (He ii 1.01m/Å) 2.5)|
|WN||He ii||log (He i 1.08)/||log (P/||(He i 1.70/||(Br/||Notes||Templates|
|Subtype||1.012m||He ii 1.01)||He ii 1.16)||He ii 1.69)||He ii 2.19)|
|7||3300||0.2||–0.3 (o)||+0.4:||–0.2 (o)||He ii 2.19 He i+N iii 2.11||WR77sc|
|6||2600||+00.2||–0.3 (o); –0.3 (h)||–0.3:||–0.2 (o); –0.2 (h)||He ii 2.19 He i+N iii 2.11 N v 2.10||WR75, WR134|
|4||2400||–0.50.5||–0.3 (o)||–0.5:||–0.2 (o); –0.2 (h)||He ii 2.19 He i+N iii 2.11 N v 2.10||WR6, WR18|
|2–3||2550||–1||–0.5 (o); –0.5 (h)||–1.0||–0.4 (o); –0.4 (h)||He ii 2.19 N v 2.10 He i+N iii 2.11||WR2, BAT51|
|WN/C||(C iii 0.97/||(C iv 1.19/||(C iv 1.74/||(C iv 2.07/||(C iv 2.43/||Notes||Templates|
|Subtype||He ii 1.01)||He ii 1.16)||He i-ii 1.7)||He i+C iii 2.11)||He ii 2.34)|
|WC||He ii||(He i 1.08/||(C iii 1.20/||(He i-ii 1.7/||(C iii 2.11/||(C ii 0.99/||Notes||Templates|
|Subtype||1.190m||He ii 1.01)||C iv 1.19)||C iv 1.74)||C iv 2.07)||C iii 0.97)|
|9||850||+1.1||+0.6||0.3||+0.10.3||–1.10.2||C ii 1.78 C iv 1.74||WR92, WR103|
|4||3300||–0.1||–0.60.1||–1.20.2||–0.7||–1.5||C ii 0.99 absent||WR143,BAT11|
|WO||C iv||(O vi 1.07/||(He ii 1.16/||(O vi 1.46+He ii 1.47)/||(C iv-iii 2.07-2.11/||(C iii 0.97/||Notes||Templates|
|Subtype||1.74m||C iv 1.19)||C iv 1.19)||C iv 1.74)||(C iv 2.43 + O vi 2.46)||He ii 1.01)|
|4||3600||–0.8||–0.7||–0.3||0.2||–0.7||O vi 1.075, 1.46, 2.46||LH41-1042|
|3||4200||–0.8||–0.7||–0.5||0.0||C iii weak||O vi 1.075, 1.46, 2.46||WR93b|
|2||6300||–0.8||–1||–0.5||–1.0||C iii absent||O vi 1.075, 1.46, 2.46||WR102|
|WR||WN||He ii 1.01||He i 1.08||P||N v 1.11||He ii 1.16||P||He ii 1.48||N v 1.55||He ii 1.69||He i 1.70||He i 2.06||N iii-v 2.11||Br||He ii 2.19||Note||Data|
|Note: 2.10 = N v 2.100; 2.11 = He i 2.112 + N iii 2.116; 2.19 = He ii 2.189|
|WR||WC||C iii 0.97||C ii 0.99||He ii 1.01||He i 1.08||He ii 1.16||C iv 1.19||C iii 1.20||C iv 1.43||He i-ii 1.70||C iv 1.74||C ii 1.78||He i 2.06||C iv 2.07||C iii 2.11||Data|
|FWHM||(He i 1.08/||(P/||(He i 1.70/||(Br/||(He i+N iii 2.11/||SpT||SpT|
|km s||He ii 1.01)||He ii 1.16)||He ii 1.69)||He ii 2.19)||N v 2.10)|
|FWHM||(He i 1.08/||(C iii 1.20/||(He i-ii 1.70/||(C iii 2.11/||(C ii 0.99/||SpT||SpT|
|km s||He ii 1.01)||C iv 1.19)||C iv 1.74)||C iv 2.07)||C iii 0.97)|
|S99 (Shara et al., 1999), H05 (Hopewell et al., 2005), M13 (Miszalski et al., 2013), K15 (Kanarek et al., 2015) (He i 2.112 + N iii 2.116/He ii 2.189)|
5 New Galactic Wolf-Rayet stars
We have identified 16 new Wolf-Rayet stars, which we have assigned Galactic WR numbers, in accordance with the current IAU convention (see Appendix of RC15). Here we discuss their spectral types, spatial location and their potential association with Scutum-Crux or other spiral arms. Near-IR spectra of the new WR stars are presented in Figures 2 (IJ) and 3 (HK), together with our NTT/SOFI observations of WR75aa, WR75c, the recently discovered WN star WR75-30 (Kanarek et al., 2015), plus previously unpublished NTT/SOFI spectroscopy of WR62a, WR68a, WR93a, WR107a, as discussed above. line measurements are provided in Table 4 and 5 for WN and WC stars, respectively, with diagnostic line ratios presented in Table 6.
5.1 Classification of the new WR stars
5.1.1 Broad-lined WN stars
Only two of the new WN stars, WR64–6 and WR70–14, are identified as broad-lined WN stars, owing to their He ii 1.01m line widths (FWHM 1900 km s) and strengths (/Å 2.4), albeit WR64-6 only narrowly complies with the second criterion. A WN6b subtype is favoured for WR64–6 from its He i 1.08/He ii 1.01m ratio which is supported by (N iii + He i 2.11) N v 2.10, while both hydrogen criteria (involving P and Br) indicate no hydrogen, so we adopt WN6b for this star. For WR70–14, the He i 1.08/He ii 1.01m ratio is somewhat ambiguous, consistent with WN4–6b, but (N iii + He i 2.11) N v 2.10 favours WN4b. This is supported by weak N v 1.11m in the J-band (Fig. 2). Again, there is no evidence for atmospheric hydrogen from our criteria (Fig. A4, available online), so WN4b is assigned to WR70–14.
5.1.2 Narrow-lined WN stars
The remaining 11 WN stars are a relatively homogeneous group, almost all classified as either WN5o, WN6o or WN7o stars, with only WR47–5 showing evidence of hydrogen so we consider these according to their subtype.
The two highest ionization narrow-lined stars are WR56-1 and WR70-15, according to their He i 1.08/He ii 1.01 ratios (Fig. A3, available online), which indicate WN5 for both stars. This is supported by the He i 1.70/He ii 1.69 ratio for WR56-1, although an earlier WN3–4 subtype is favoured by He i 1.70/He ii 1.69 for WR70-15. We also consider their K-band morphologies, for which (N iii + He i 2.11) N v 2.10 in both cases, indicating WN5–6. Neither star shows any evidence for atmospheric hydrogen from Fig. A4 so we adopt WN5o for both stars.
The majority of our narrow-lined WN stars are WN6 stars according to their He i 1.08/He ii 1.01 ratios (Fig. A3), with He i 1.70/He ii 1.69 suggesting WN4, 5 or 6. As with the WN5o stars considered above, we also consider the K-band morphology, for which either (N iii + He i 2.11) N v 2.10, implying WN5–6 or (N iii + He i 2.11) N v 2.10, implying WN6. Only WR47–5 indicates the presence of (modest) hydrogen from Fig. A4, such that we classify it as WN6(h), but favour WN6o for WR60-8, WR64-2, -3, -4, -5 and WR72-5.
Of the remaining stars only WR75-31 was observed in the IJ-band with SOFI, for which He i 1.08/He ii 1.01 indicates a WN7–8 subtype, with He i 1.70/He ii 1.69 also providing an ambiguous WN7–8 classification. Its K-band morphology strongly favours WN7 since (N iii + He i 2.11) He ii 2.19, while there is no evidence for atmospheric hydrogen in WR75-31 from Fig. A4, such that we assign WN7o to this star. WR76-11 was observed solely in the H and K bands, but closely resembles WR75-31 such that we classify it as WN7o, as with WR75-30. None of the new WN stars qualify as WN/C stars, since C iii 0.971m, C iv 1.19m, 1.74m, 2.07m are weak/absent.
5.1.3 WC stars
Three of the new WR stars are carbon sequence WC stars. Considering the primary diagnostics for WR46–16, WC7 is favoured from the C iv 1.19/C iii 1.20 ratio, WC4–6 from the C iii 0.97/C ii 0.99 ratio, and WC5–7 from the C iv 2.07/C iii 2.11 ratio. Secondary indicators suggest WC5–7 from He i 1.08/He ii 1.01, WC5–6 from C iii 0.97/He ii 1.01 and WC6–7 from He i-ii 1.7/C iv 1.74. Overall, we adopt WC6–7 reflecting the tension in primary indicators for WR46–16 (Fig. A7, available online).
WR60–7 is classified as WC8, WC7, WC7–8 from primary diagnostics C iv 1.19/C iii 1.20, C iv 2.07/C iii 2.11 and C iii 0.97/C ii 0.99, respectively. Secondary criteria C iii 0.97/He ii 1.01 and He i-ii 1.7/C iv 1.74 indicate WC6–7, while He i 1.08/He ii 1.01 favours WC7–8. Overall, WC7–8 is selected for WR60–7, reflecting the lack of a consensus amongst primary criteria (Fig. A7).
Finally, primary diagnostics C iv 1.19/C iii 1.20, C iv 2.07/C iii 2.11 and C iii 0.97/C ii 0.99, imply WC8–9, WC8 and WC8–9 for WR70–13, while C iv 1.74 C ii 1.78 indicates WC8. Consequently we adopt WC8 for WR70–13, which is supported by our secondary indicator He i-ii 1.7/C iv 1.74, with He i 1.08/He ii 1.01 and C iii 0.97/He ii 1.01 consistent with either WC8 or WC9 (Fig. A7).
Approximately 40% of the Galactic WR population are observed in multiple systems (van der Hucht, 2001). This is a lower limit on the true binary fraction, since no systematic survey has been carried out. It is therefore highly likely that some of the newly discovered WR stars are in fact multiple systems. Direct detection of companion stars, usually main-sequence OB stars, is not possible with the current dataset since their absorption lines are generally weak with respect to the strong WR emission lines.
It is, however, possible to infer the presence of a companion star by considering the equivalent width of near-IR emission lines which will be diluted by the continuum of a companion star, and/or dust for the case of some WC+OB systems since dust formation is an indicator of binarity in WC stars.
Since a companion star and/or thermal dust emission will not reduce line widths, a weak line compared to single stars at a specific FWHM is suggestive of binarity. In Figure 4 we compare the FWHM (km s) and equivalent widths (in Å) of strong, isolated lines in apparently single Galactic WN stars (He ii 1.012µm) and WC stars (C iv 1.736µm) with newly discovered WR stars. We also include weak-lined WN stars with intrinsic absorption lines (WR24, WR87, WR108) which could be mistaken for WN+OB stars, plus dusty WC stars (WR121, WR75aa), whose near-IR emission lines are diluted by hot dust.
Of the newly identified stars, the majority of WR stars possess emission line strengths which are charactistic of single stars. From Fig. 4(a) two exceptions are WR75-31 (WN7(h)) and WR64-4 (WN6o) which possess weak emission for their He ii 1.012m FWHM. Both are potential binaries, although WR64-4 is the strongest candidate, such that we revise its spectral type to WN6o+OB. In contrast, WR75-31 has an overall relatively strong emission line spectrum, albeit with an anomalously weak (and low S/N) He ii 1.0124m line.
Of the WC stars, none possess unusually weak emission lines based on their C iv 1.736m FWHM (Fig. 4(b)). However, the increased dilution of WC emission lines from 1m to 2.5m arising from hot dust in WCd systems also severely modifies equivalent width ratios of C iii-iv lines. For example, (C iii 2.11)/(C iii 0.97) = 0.3 for WR88 (WC9) but hot dust in WR121 (WC9d) reduces this ratio to 0.05. A similar reduction in line strength is observed for prominent He ii lines, with (He ii 2.19)/(He ii 1.28) = 0.5 for WR88 and 0.17 for WR121. WR135 is a prototypical non-dusty WC8 star with (C iii 2.11)/(C iii 0.97) = 0.2, with a ratio of 0.2 for WR60-7 but only 0.1 for WR70-13, suggestive of dust dilution in the latter. Indeed, WR60-7 (WC7–8) and WR70-13 (WC8) possess similar J-H colours, yet the latter has 0.5 mag higher K– colour (Table 1), so we amend its spectral type to WC8d. Indeed, WR70-13 is offset from the other WC stars in the reddening free Q1 vs Q2 comparison in Fig. 1. Turning to WR46-18, (C iii 2.11)/(C iii 0.97) 0.3 for non-dusty WC6–7 stars, with a ratio of 0.4 for WR46-18, arguing against dust emission in this instance.
Discovery of a hard X-ray source associated with any of the WR stars would be highly indicative of stellar wind collision in a massive binary. Indeed, several WR stars towards the Galactic Centre, coincide with hard X-ray sources (e.g., Mauerhan et al. 2010, Nebot Gómez-Morán et al. 2015). From a search of the XMM Newton science archive and the Chandra source catalogue (1.0), fields including WR60-7 (WC7–8), WR60-8 (WN6o), and WR64-4 (WN6o) had been imaged by Chandra ACIS, although none revealed a source at the location of the WR star.
5.3 Spatial location of the new WR stars
We have estimated distances to the new WR stars by adopting an absolute K-band magnitude based on the assigned spectral subtype. To do this, we followed the approach of RC15, which we briefly summarise here. To calculate the foreground dust extinction to each new WR star, we used subtype-specific intrinsic JH and HK colours to measure a colour excess. Using the near-IR extinction law of Stead & Hoare (2009), we thereby obtained two measures of extinction in the K-band, of which an average was taken. Distances were then calculated using the standard relation between absolute and apparent magnitude. We used the uncertainties on calibrated absolute magnitudes, given by RC15, to calculate upper and lower bounds on the distances calculated. In Table 7 we provide interstellar extinctions, distance moduli/distances, with uncertainties, for every new WR star. Typical extinctions are 0.2 mag, with characteristic distances of 9.73.8 kpc.
This method inherently assumes the WR star is the sole (or dominant) contributor of near-IR flux to each source. Recalling Sect. 5.2, this assumption is justified for new WN stars with the exception of WR64-4, the emission line strengths of which suggest a significant contribution from a companion source, implying a larger distance. In addition, we provide a second distance estimate to WR70-13 in Table 7 since there is evidence for a contribution by circumstellar dust to the K-band. Adopting the absolute magnitude of a WC8 star which is dominated by hot dust would significantly increase the distance to WR70-13 from 5.3 to 10.7 kpc, though in reality the dust contribution is likely to be modest such that an intermediate distance is more realistic.
5.4 Association of WR stars with the Scutum-Crux and other spiral arms
In Figure 5 we present the locations of the new WR stars on a top-down view of the Galactic disk, together with WR stars mapped by RC15 (including approximately half of the currently known population). Over-plotted are the locations of the three main spiral features detected in the 4th Galactic quadrant (), assuming kpc. We assume each arm is a logarithmic spiral, parameterised as:
in which is the Galactocentric radius of the observed tangent, is the angle measured anti-clockwise about the origin from the positive x-axis, and is the angle at which the observed tangent is located. The parameter is the pitch angle of the arm. The calculation of requires a measurement of the longitude of the tangent to each arm (), and the Galactocentric radius of the Sun ().
Vallée (2014) catalogued the observed tangents to spiral arms in the Galaxy, and calculate averages of measurements using different tracers. Subsequently, Vallée (2015) use these observations to measure pitch angles for individual spiral arms. From these studies we adopt for the Sagittarius-Carina arm, for the Scutum-Crux arm, and for the Norma (3kpc) arm.
The new WR stars appear to be evenly distributed throughout the section of Galactic disk observed; neither they, nor the previously mapped WR stars, show any obvious association to the spiral features in the region. Indeed, no pattern resembling a spiral arm can be seen in the distribution of WR stars, as would be expected if they were tightly confined to spiral arms, albeit affected by a systematic offset in distance measurements. The upcoming second data release (DR2) from Gaia should address this general question, although the majority of the newly discovered WR stars are too faint for reliable parallaxes with Gaia. Indeed, only 7 of the 17 stars are included in the Gaia first data release (DR1 Gaia Collaboration et al., 2016), with G = 17.4 – 20.1 mag.
|Cluster||Alias||Ref||WN (, arcmin)||WC (, arcmin)||Ref|
|VVV CL 011||298.506||–0.170||5.0:||0.1||d, l||WR46-17 (0.0)||d|
|Mercer 30||298.756||–0.408||7.2||0.3||a, j||WR46-3 (0.2), WR46-4 (0.1)||a|
|WR46-5 (0.1), WR46-6 (0.2)||a|
|C 1240-628||Hogg 15||302.047||–0.242||3.0||3.5||k||WR47 (1.6)||b|
|C 1309-624||Danks 1||305.339||+0.080||4.2||0.75||c, k||WR48-8 (0.6), WR48-9 (0.6)||WR48a (1.3), WR48-3 (1.9)||c|
|WR48-10 (0.6), WR48-6 (2.7)||WR48-4 (2.4)||c|
|C 1310-624||Danks 2||305.393||+0.088||4.2||0.75||c, k||WR48-2 (0.6)||c|
|VVV CL 036||312.124||+0.212||2.0:||0.8||d, l||WR60-6 (0.1)||d|
|VVV CL 041||317.109||+0.281||4.2||0.5||e, l||WR62-2 (0.2)||e|
|Pismis 20||320.516||–1.200||3.6||2.0||k||WR67 (2.1)||i|
|Mercer 70||329.697||+0.584||7.0||0.4||f, j||WR70-12 (0.4)||f|
|VVV CL 073||335.894||+0.133||4.0:||0.3||d, l||WR75-25 (0.1), WR75-26 (0.1)||d|
|VVV CL 074||336.373||+0.194||6.0:||0.55||d, l||WR75-28 (0.1), WR75-29 (0.0)||WR75-27 (0.3)||d|
|Mercer 81||338.384||+0.111||11.0||0.6||g, j||WR76-2 (0.6), WR76-3 (0.7)||g|
|WR76-4 (0.9), WR76-5 (0.8)||g|
|WR76-6 (0.9), WR76-7 (0.9)||g|
|WR76-8 (0.9), WR76-9 (1.0)||g|
|C 1644-457||Westerlund 1||339.555||–0.399||3.8||1.2||m, n, k||WR77a (1.7), WR77c (0.8)||WR77aa (4.1), WR77b (4.8)||h|
|WR77d (1.1), WR77e (0.4)||WR77g (0.1), WR77i (0.9)||h|
|WR77f (0.4), WR77h (0.1)||WR77l (0.6), WR77m (0.4)||h|
|WR77j (0.7), WR77k (0.4)||WR77n (1.7), WR77p (0.8)||h|
|WR77o (0.4), WR77q (0.5)||h|
|WR77r (0.8), WR77s (0.5)||h|
|WR77sa (0.4), WR77sb (2.0)||h|
|WR77sc (0.8), WR77sd (2.8)||h|
|a: Kurtev et al. (2007); b: Sagar et al. (2001); c: Davies et al. (2012b); d: Chené et al. (2013); e: Chené et al. (2015); f: de la Fuente et al. (2015); g: Davies et al. (2012a); h: Crowther et al. (2006) i: Lundstrom & Stenholm (1984) j: Mercer et al. (2005) k: Dias et al. (2002) l: Borissova et al. (2011) m: Kothes & Dougherty (2007) n: Koumpia & Bonanos (2012)|
|WISE||Ref||Cluster?||WN (, arcmin)||WC (, arcmin)||Ref|
|G298.224–0.334||11.1||5.0||a||WR46-7 (5), WR46-18 (8)||b, p|
|G298.529–0.251||10.5||16.2||a||VVV CL 011||WR46-8 (15.6), WR46-9 (3.5), WR46-17 (4.9)||WR46-10 (7.6)||b, c, d, e|
|WR46-15 (16.9), WR46-2 (22.4)||g|
|G298.862–0.432||10.7||4.8||a||Mercer 30||WR46-3 (6.6), WR46-4 (6.6), WR46-5 (6.6),||f|
|G305.233+0.110||4.9||11.3||a||Danks 1/2||WR48-6 (5.5), WR48-10 (6.5), WR48-7 (7.7)||WR48-1 (8.6), WR48-3 (6.0)||e, g, h, i|
|WR48-9 (6.8), WR48-8 (6.8)||WR48-4 (8.2), WR48a (8.4)||h, i|
|G311.991+0.219||6.0||a||VVV CL 036||WR60-6 (7.8)||c|
|G317.030+0.028||24.4||a||VVV CL 041||WR62-2 (15.7), WR62-1 (19.4), WR63 (34)||h, j|
|G335.794+0.153||13.7||a||VVV CL 073||WR75-6 (19.7), WR75-25 (6.1), WR75-26 (6.1)||WR75b (16.1), WR75-15 (12.4)||h|
|G336.446–0.198||10.4||a||WR75-9 (11.7), WR75-21 (10.3)||n|
|G338.430+0.048||5.8||a||Mercer 81||WR76-2 (3.5), WR76-3 (3.5), WR76-4 (3.5)||o|
|WR76-5 (3.5), WR76-6 (3.5), WR76-7 (3.5)||o|
|WR76-8 (3.5), WR76-9 (3.5)||o|
|a: Anderson et al. (2014) b: Mauerhan et al. (2009) c: Chené et al. (2013) d: Shara et al. (2009) e: Hadfield et al. (2007) f: Kurtev et al. (2007) g: Mauerhan et al. (2011) h: van der Hucht (2001) i: Davies et al. (2012b); j: Chené et al. (2015) k: Marston et al. (2013) l: Roman-Lopes (2011) m: Wachter et al. (2010) n: Shara et al. (2012) o: Davies et al. (2012a); p: This work : Alternatively WR62-2 and WR62-1 may be associated with WISE H ii region G317.236+0.516|
6 Spatial distribution of Wolf-Rayet stars and other massive stars in the Milky Way
6.1 Association of WR stars with star clusters?
If the majority of WR progenitors are born in relatively massive star-forming regions, one might expect them to be in close proximity to their natal star cluster. We have compared the spatial location of the new WR stars with young star clusters from Dutra et al. (2003), Mercer et al. (2005) and Borissova et al. (2011). None of the new WR stars are located within known clusters. At face value, this might be considered to be surprising, but Lundstrom & Stenholm (1984) in the most extensive study of the environment of Galactic WR stars to date, found that only 10–30% of WR stars are located in star clusters.
Since then, the known WR population has quadrupled, so we provide revised statistics for the 298 survey region as a whole, involving 190 + 16 = 206 Galactic Wolf-Rayet stars, comprising 119 WN, 1 WN/C and 86 WC stars. Lundstrom & Stenholm (1984) considered a WR star to be associated with a star cluster if its projected distance was within two cluster radii. Here, we soften this requirement, extending the potential association to 4 cluster radii, utilising published cluster centres and radii from Dias et al. (2002), Dutra et al. (2003), Mercer et al. (2005), Borissova et al. (2011). Overall, 55 WR stars (27%) are associated with a total of 12 star clusters, as shown in Table 8, although it is notable that 44% of all the WR cluster members in our survey range are located within a single open cluster, Westerlund 1 (C06). Consequently 73% of WR stars in the survey region are not associated with a star cluster.
If the majority of WR progenitors are born in relatively massive star forming regions, why are so few currently associated with clusters? The lower WR mass limit for single rotating stars at Solar metallicity is 22 star (Meynet & Maeder, 2003). Empirically, such stars are observed in clusters with (Larson, 2003; Weidner et al., 2010). Stochasticity in the sampling of the initial mass function will result in some massive stars originating in low mass (10) clusters/star-forming regions, as demonstrated theoretically by Parker & Goodwin (2007). Therefore, some WR stars could be associated with low mass clusters, with other members of the star forming region inconspicuous. Indeed, Vel (WC8+O), the closest Galactic WR star, is located in a very low mass star-forming region (Jeffries et al., 2014). Other members of such star-forming region would be very difficult to identify at the typical distance of Galactic WR stars.
There is evidence that some dense, young massive clusters rapidly achieve virial equilibrium (e.g. Hénault-Brunet et al., 2012), such that they would retain the bulk of their stellar content over the representative WR ages of 5–10 Myr (Meynet & Maeder, 2003). Indeed, the bulk of the WR stars associated with Westerlund 1 lie within the 1.2 arcmin radius – corresponding to 1.4 pc at a distance of 3.8 kpc – such that our 4 cluster radius limit is equivalent to 5.5 pc. Of course not all open clusters are dense or bound. For example, Hogg 15 is a sparsely populated cluster with a much larger radius of 3 pc, at its distance of 3 kpc, so will be in the process of dispersing. Consequently the general lack of an association with star clusters is not wholly unsurprising. Indeed, inspection of the Galactic O Star Catalogue (GOSC v3.0 Maíz Apellániz et al., 2013) reveals 50 (optically visible) O stars in our survey region. Of these, only 18 (36%) are associated with a star cluster, not significantly larger than WR stars since the bulk of these are late-type O stars with comparable ages to many WR stars.
Of course, not all massive stars originate from star clusters. Bressert et al. (2010) demonstrated that only a small fraction of star formation in the Solar Neighbourhood originates from dense environments. It is probable that the majority of OB stars originate from intermediate density regions, i.e. OB asociations and/or extended star-forming regions (see Parker & Dale, 2017). Indeed, Wright et al. (2016) have unambiguously demonstrated that the distributed massive stars in Cyg OB2 did not originate from a dense star cluster.
Lundstrom & Stenholm (1984) found that 50% of optically visible WR stars were located in OB associations. It is not possible to calculate the fraction of IR-selected WR stars that lie within OB associations since the latter are limited to nearby optical surveys. Instead, it is necessary to compare the location of Wolf-Rayet stars with infrared catalogues of star-forming regions. We have compared the locations of all 206 WR stars in our survey region with confirmed H ii regions from the all-sky WISE catalogue of Galactic H ii regions from Anderson et al. (2014). In total, 53 stars are located within 1.5 of the H ii region, representing 26% of the total WR population, as shown in Table 9. Of course, a subset of these WR stars will be foreground sources, so the quoted statistics represent strict upper limits.
The majority of these WR stars are associated with complexes at G298 (Dragonfish nebula), G305 and G338. By way of example, as many as twelve WR stars are associated with G298 ‘Dragonfish’ complex (Rahman et al., 2011). The VVV CL011 and Mercer 30 clusters are coincident with this region, although radio distances of 10.5 kpc significantly exceed spectrophotometric cluster distances. Similar issues affect the association of Mercer 81 with G338.430+0.048, and VVV CL 041 with G317.030+0.028. Finally, stellar winds and/or supernovae associated with Westerlund 1 have sculpted an IR cavity, such that there is no longer an associated H ii region.
Overall, it is apparent that WR stars are rarely located within known H ii regions, albeit with some notable exceptions which includes the Carina Nebula in the Milky Way. Typical radii of star-forming regions are 9 arcmin, so for representative radio-derived distances of 7.5 kpc to star forming regions, we include stars within a projected distance of 30 parsec from the centre of the H ii region. Again, it is possible that stars have migrated away from where they formed. Typical velocity dispersions of stars in such regions are several km/s, corresponding to several pc/Myr, such that the WR progenitor could have travelled several ten’s of parsec. It is important to stress that the parent star-forming region of a WR star would not necessarily display significant radio free-free emission after 5–10 Myr (Meynet & Maeder, 2003) since most O stars will have died if there had been a single burst of star formation. Indeed, only 23 (46%) of the 50 GOSC optically visible O stars in our survey region lie in an OB association (Cen OB1, Nor OB1, Ara OB1).
Regardless of whether WR progenitors form in clusters or lower density environments, there are other explanations for their relative isolation. To have travelled more than a few 10s of parsec from their birth site, the WR progenitor may have been ejected via dynamical effects or following a supernova kick from a close companion. The former, involving 3-body dynamical interactions, is favoured in dense stellar environments, in which the fraction of massive stars ejected is inversely proportional to the stellar mass of the cluster (Fujii & Portegies Zwart, 2011). Therefore, the population of WR stars dynamically ejected in this way will be dominated by those from low to intermediate mass clusters, explaining why the 10 cluster Westerlund 1 has successfully retained the majority of its WR population. Historically, a supernova origin for runaway WR stars has been considered to be a major factor affecting their spatial distribution (Eldridge et al., 2011). This requires a close binary companion, and assumes that all massive stars whose initial masses exceed undergo a core-collapse SN. Of course, such stars may collapse directly to a black hole, or form a black hole via fallback, so their companion would not necessarily receive a significant supernova kick (e.g. O’Connor & Ott, 2011). Therefore, only a small fraction of isolated WR stars likely originate in this way.
Overall, the most promising scenario for the low observed frequency of WR stars associated with star clusters is via dynamical ejection, but this alone doesn’t explain the very high fraction of WR stars in the field. Instead, it is likely that a majority of Galactic WR progenitors do not form within dense, high mass star clusters. The observed high fraction of WR stars located in the Galactic field can therefore be attributed to a combination of dynamical ejection from star clusters, and their origins in modest star-forming regions which subsequent dissolve into the field. The latter will not necessarily be recognisable as an OB association during the WR stage since the majority of O stars will have ended their lives after 5–10 Myr, with similar arguments applying to isolated H ii regions. Some distant WR stars will not be recognised as being associated with a star forming region if other stars in the region possess low masses, as is the case for Vel (Jeffries et al., 2014).
6.2 Relative isolation of WR stars and Luminous Blue Variables
The general lack of association between Wolf-Rayet stars and O stars – except for the minority located in young, dense star clusters (e.g. NGC 3603, Westerlund 1) – is relevant to the ongoing debate about the nature of Luminous Blue Variables (Humphreys et al., 2016; Smith, 2016). Historically LBVs were considered to be massive stars in transition towards the WR stage. In part, this association arose from the spectroscopic evolution of LBVs to the WN phase (e.g. AG Car Smith et al., 1994) and LBV outbursts from WN stars (e.g. R127, HDE 269582, Walborn et al., 2017). Smith & Tombleson (2015) have challenged this view, finding that LBVs in the Milky Way and LMC are more isolated (from O stars) than Wolf-Rayet stars. They have argued that they possess lower masses, with their high luminosities arising from being the mass gainers (former secondaries) within close binary systems.
Their conclusions were largely based upon the spatial distribution of O stars, WR stars and LBVs in the LMC, owing to visual studies of Galactic massive stars being severely restricted by dust extinction. Still, the reliance on SIMBAD for catalogues of O stars hinders their conclusions owing to severe incompleteness for both galaxies, while Humphreys et al. (2016) argued for a mass separation between high luminosity (classical) and low luminosity LBVs (though see Smith, 2016).
Kniazev et al. (2015, 2016) provide the current census of 17 confirmed LBVs in the Milky Way, which is restricted to confirmed spectroscopic variables. Their criteria therefore exclude candidate LBVs possessing ejecta nebulae, which Bohannan (1997) had previously argued should be an additional discriminator. Only three spectroscopically variable LBVs lie within our Galactic survey region – WS1 (Kniazev et al., 2015), Wray 16-137 (Gvaramadze et al., 2015) and Wd1-W243 (Clark et al., 2005) – so we need to consider the global Milky Way population of LBVs and WR stars.
Of the known WR content in the Milky Way, 27% are members of star clusters (Crowther, 2015). Meanwhile, 4 of the 17 spectroscopically variable LBVs are located in star clusters – W243 in Westerlund 1, Car in Trumpler 16, GCIRS 34W in the Galactic Centre cluster and qF 362 in the Quintuplet cluster (Geballe et al., 2000)) – comprising 24% of the total, so their global statistics are comparable. Indeed, a number of widely accepted LBVs known to be associated with star clusters are omitted from the compilation of Kniazev et al. (2015). These include the Pistol star (qF 134), another member of the Quintuplet cluster (Figer et al., 1998), and the LBV in the 1806-20 cluster (Eikenberry et al., 2004; Figer et al., 2004).
5 of the 17 spectroscopically variable LBVs are potentially associated with star-forming regions (Anderson et al., 2014) from a similar exercise to that discussed above for WR stars, namely Car, Wray 16-137, HD 168607, AFGL 2298 and G24.73+0.69, namely 29% of the total. Consequently, the overwhelming majority of WR stars and LBVs are located in the Galactic field, away from star clusters or star-forming regions. Overall, there is no significant difference in the spatial distribution of WR and LBVs in our Galaxy, with a quarter of both populations associated with young massive star clusters.
Smith & Tombleson (2015) proposed that LBVs generally arise from significantly lower mass progenitors than WR stars, challenging the hitherto scenario that LBVs are transition objects between the O and WR phases. However, those star clusters which host LBVs are relatively young (4–6 Myr Clark et al., 2005; Bibby et al., 2008; Liermann et al., 2012; Schneider et al., 2014), putting aside Car666 Car is unusual since it is associated with an even younger star cluster, Trumpler 16, which also hosts the massive main-sequence O2.5 If/WN6 star WR25 (Massey & Johnson, 1993; Crowther & Walborn, 2011). Not only are the statistics of WR and LBV association with star clusters comparable, but crucially 4 young Milky Way clusters hosting LBVs – Westerlund 1, CL 1806-20, the Galactic Centre and the Quintuplet – also contain (classical) WR stars.
Smith & Tombleson (2015) argued that LBVs arise from mass gainers in close binary systems (Langer, 2012; de Mink et al., 2014). Mass gainers will be rejuvenated, yielding an apparently younger star than the rest of the population (Schneider et al., 2014). The presence of spectroscopically variable LBVs (Wd1-W243, GCIRS 34W, qF 362) plus leading LBV candidates (qF 134, LBV 1806-20) in young clusters with progenitor masses in the range 30–40 and coexisting with O stars and WR stars, should permit their scenario to be tested. Of course, LBVs in such systems might be the mass-gaining secondaries whose primaries have already undergone core-collapse. However, LBV 1806-20 is an SB2 system, whose companion is not a WR star since its exhibits He i absorption lines. This appears to rule out the Smith & Tombleson (2015) scenario in this instance.
If LBVs are rejuvenated secondaries in close binaries spanning a wide range of masses, they should also be present in older, massive star clusters such as those hosting large red supergiant (RSG) populations. However, Smith & Tombleson (2015) report that LBVs and RSGs do not share a common parent population, and LBVs are not known within RSG-rich clusters at the base of the Scutum-Crux arm (RSGC 1–3, Figer et al., 2006; Davies et al., 2007; Clark et al., 2009). The presence of 50 RSG in these clusters and absence of LBVs argues either argues against the Smith & Tombleson (2015) scenario, or requires a short ( yr) lifetime for such LBV, in view of the 1 Myr RSG lifetime. The only potential LBV within these RSG rich clusters identified to date is IRAS 18367—0556 in RSGC 2, which possesses a mid-IR ejecta nebula (Davies et al., 2007). Overall, it is apparent that LBVs span a wide range of physical properties (Smith et al., 2004), though the same is true for WR stars, some of which are located in significantly older star clusters (recall Table 9 from C06), albeit solely Westerlund 1 hosts RSG and WR stars.
We have undertaken a near-IR spectroscopic search for Wolf-Rayet stars along the Scutum-Crux spiral arm, based upon 2MASS and GLIMPSE photometric criteria (Mauerhan et al., 2011; Faherty et al., 2014). Observations of 127 candidate stars ( mag) resulted in the confirmation of 17 WR stars (14 WN, 3 WC), representing a success rate of 13%, comparable to previous IR-selected studies based on similar criteria (Hadfield et al., 2007; Mauerhan et al., 2011). More sophisticated techniques are clearly required for optimising future spectroscopic searches. As we have found, large numbers of other stellar types (young stellar objects, Be stars) lie in the same location as WR stars within individual colour-colour diagrams, but may not do so in a multi-dimensional parameter space. Therefore, future progress might entail a Bayesian approach to optimising searches for candidate WR stars.
We have extended our earlier near-IR classification scheme (C06) to cover all YJHK bands and all subtypes, including WN/C and WO subtypes. This has been tested on several recently discovered WR stars for which optical classifications have previously been obtained. In general, the near-IR criteria are successful in obtaining reliable subtypes, including the presence of atmospheric hydrogen in WN stars, and identifying transition WN/C stars. However, inferences are weaker if limited wavebands are available, the spectral resolution is modest, and/or the signal-to-noise obtained is low.
The majority of newly discovered WR stars are weak-lined WN5–7 stars, with two broad-lined WN4–6 stars and three WC6–8 stars. Therefore, despite the low success rate, our goal of extending the spectroscopic confirmation of WR stars to the far side of the Milky Way has been successful. Based on the absolute magnitude calibration of C15, inferred distances (10 kpc) extend beyond previous spectroscopic surveys, with mag. Spectroscopic searches beyond the Galactic Centre ( mag) will be significantly more challenging.
Only a quarter of WR stars in the selected Galactic longitude range are associated with star clusters and/or H ii regions. We suggest that this arises from a combination of dynamical ejection from (modest stellar mass) clusters and formation within lower density environments (OB associations/extended H ii regions). We also revisit the recent controversy about the association between LBVs and WR stars, or lack thereof. Considering the whole Milky Way, 27% of WR stars are hosted by clusters, comparable to that of spectroscopically variable LBVs (4 from 17 stars). More significantly, several young clusters with main sequence turn-off masses close to 30–40 host classical WR stars and LBVs, permitting Smith & Tombleson (2015)’s suggestion that some LBVs are rejuvenated mass gainers of close binaries to be tested. Specifically, the non-WR companion to LBV 1806-20 and absence of LBVs in RSG-rich star clusters argue against this scenario.
Returning to the main focus of this study, until more robust distances to WR stars - and in turn absolute magnitudes - can be established by Gaia there are legitimate concerns about the completeness of surveys for different subtypes, especially the challenge of identifying faint, weak emission line WN stars with respect to WC stars (Massey et al., 2015a). Near-IR narrow-band photometric searches suffer from dilution of emission lines by companion stars and hot dust emission from WC binaries, while our broad-band near to mid-IR photometric approach is limited by the low spatial resolution of Spitzer. Massey et al. (2015b) have undertaken a deep optical narrow-band survey of the LMC, revealing a population of faint, weak-lined WN3 stars (which they characterize as WN3/O3 stars). Stars with these characteristics – which would usually be classified as WN3ha according to the Smith et al. (1996) scheme – comprise a negligible fraction of Galactic WR stars (e.g. WR3), a minority of the moderately metal-poor LMC WR population, and a majority of the more significantly metal deficient SMC WR population (Hainich et al., 2015).
It is anticipated that ongoing infrared searches using a combination of continuum methods and emission line characteristics will significantly improve the completeness of WR surveys in a sizeable volume of the Galaxy within the near future.
We wish to thank Nicole Homeier, Bill Vacca and Frank Tramper for providing us with near-IR spectroscopic datasets, from NTT/SOFI, IRTF/SpeX and VLT/XSHOOTER respectively, for several WR stars. We appreciate useful discussions with Simon Goodwin and Richard Parker, and useful comments from the referee which helped improve the focus of the paper. Financial support for CKR was provided by the UK Science and Technology Facilities Council. PAC would like to thank ESO for arranging emergency dental care in La Serena immediately after the 2015 observing run.
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Appendix A: Near-IR spectral classification of WR stars
Appendix B: Spectroscopic observations of IR selected candidates
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Appendix A Near-IR Spectral classification of WR stars
C06 utilised He ii 1.012µm/He i 1.083µm and C iv 1.19µm/C iii 1.20µm to classify WN and WC stars from YJ-band spectroscopy, respectively, plus He ii 2.189µm/Br and C iv 2.08µm/C iii 2.11µm from K-band spectroscopy. Subtypes for weak/narrow WN4–7 stars with a He ii 1.01µm FWHM 65Å were distinguished by the ratio of N v 2.10µm/(He i + N iii 2.11µm), while WC8–9 stars were also classified from their C iii 0.97µm/C ii 0.99µm ratio. However, separating broad-lined WN4–7 stars and WC5–7 stars from near-IR spectroscopy proved to be extremely challenging. In addition, the use of He ii 2.189µm/Br as a classification diagnostic adds a further complication owing to the sensitivity of this ratio on hydrogen content as well as ionization. Consequently, we revisit these issues in the following discussion.
a.1 WN stars
The primary classification diagnostics of WN stars involve helium (He ii 5411/He i 5876) and nitrogen (N iii 4640, N iv 4058, N v 4610) lines, with strong, broad-lined WN stars (WNb or WN-s) further distinguished from weak narrow-lined stars (WN, WN-w). In common with C06, we identify strong/broad WN stars with (He ii 5411) 40Å, and FWHM(He ii 4686)30Å (Smith et al., 1996) using the FWHM of prominent He ii lines, 1.012m and/or 2.189m, namely 1900 km s for WNb stars, plus an equivalent width threshold of 250Å for He ii 1.0124m, and/or 75Å for He ii 2.1885m.
Central to the optical classification scheme are the ratios of N iii 4636-41, N iv 4068 and N v 4603-20 lines. At near-IR wavelengths, nitrogen lines are scarse, weak, and often blended with helium lines. The most prominent nitrogen features arise from N v (1.111, 1.552, 2.100m) in early WN subtypes, although these recombination lines require high S/N spectroscopy for broad-lined stars. In addition, N iii 2.116 is prominent in late WN stars, albeit severely blended with He i 2.112–2.113m and/or C iii 2.115m for WN/C stars.
Consequently we are largely reliant upon the relative strengths of helium lines for IR classification. In the YJ-band, He ii 1.012µm/He i 1.083µm represents an excellent classification diagnosic for most WN subtypes. He i 1.0830 is blended with P (+ He ii 1.093µm) in early subtypes and broad-lined stars, while He ii 1.0124 is blended with P (+ He ii 1.0045µm) in broad-lined WN stars. Divisions between subtypes are assigned suing optically classified WN stars, although broad-lined WN stars cannot be distinguished from this diagnostic.
He ii 1.692m/He i 1.700m provides a H-band diagnostic in narrow-lined WN stars, especially for late subtypes, although blending severely hinders its role in broad-lined stars. Figure 8 compares these ratios in WN stars, indicating that they are well correlated for most narrow-lined WN subtypes.
In the K-band, the relative strengths of N iii 2.116m (+ He i 2.112-2.113) to N v 2.100m (or He ii 2.189) provide key diagnostics in weak-lined WN stars, although this is hindered at low spectral resolution and signal-to-noise since these emission lines are generally weak. Other relevant pairs of diagnostics, He ii 1.162/P and He ii 2.189/Br provide a combination of ionization and hydrogen content information, so if classifications can be established from other diagnostics, these are helpful in distinguishing hydrogen-rich from hydrogen-deficient WN stars.
a.1.1 Hydrogen content
Conti et al. (1983) and Smith et al. (1996) employed the Pickering-Balmer series to characterize the hydrogen content of WN stars, the latter introducing ’o’, ’(h)’ and ’h’ subtypes for narrow-lined WN stars based on the strength of H + He ii 4859 (8-4) with respect to adjacent Pickering series members He ii 5411 (7–4) and He ii 4541 (9–4), with ’o’ omitted for broad-lined WN stars. To date, no attempts have been made to develop similar criteria from near-IR series (e.g. Paschen, Brackett), largely because there are additional complications with respect to optical diagostics.
Firstly, leading He ii near-IR lines are blended with other He ii series (incl. -7) whose contributions need to be taken into account (e.g. 1.488m 14–7 with 1.476m 9–6). Secondly, hydrogenic He i emission contributes significantly to the hydrogen lines (e.g. He i 5–3 at P). The former issue can be overcome fairly easily, since medium resolution spectroscopy of late WN stars WR120 and WR123 (dataset U6 from Table 2) reveals (He ii 1.167)/(1.163) 0.35 and (He ii 1.488)/(1.476) 0.35. However, contributions from neutral helium lines to hydrogen features are more problematic. By way of example, one would expect the observed equivalent width of 1.28m to be 10Å for the hydrogen-free WN8 star WR123 (Crowther et al., 1995b) if He ii (10-6) is the dominant contributor, since (He ii 11-6) = 9.40.2Å and (He ii 9-6) = 11.50.6Å, whereas (1.281m He ii + P + He i 5-3) = 701Å. Similar issues affect P and its adjacent He ii lines, 1.042m (13-6) and 1.167m (11-6). Therefore, it is not straightforward to determine the hydrogen content in low ionization WN stars from the Paschen series and their adjacent He ii lines.
Instead, we compare the equivalent width ratios of He i 1.08m/He ii 1.01m (ionization) to P/He ii 1.16m (ionization and hydrogen content) for narrow-lined WN stars in Fig. 9(a), and broad-lined WN stars in Fig. 9(b). WN stars with optical (h) or ’h’ classifications are indicated by cross (narrow-lined) or plus (broad-lined) symbols. It is apparent that WN stars with hydrogen exhibit excess P strengths at a given ionization, such that we have overlaid a pale shaded region for those stars with hydrogen classifications. A darker shaded region, vertically offset by 0.3 dex corresponds to stars with significant hydrogen. Stars close to the boundary include WR66 (WN8(h)) for which Hamann et al. (2006) have obtained 5% atmospheric hydrogen. We have adopted the same criteria for broad-lined WN stars, for which WR136, the only Galactic WNb star with some atmospheric hydrogen in our sample (Crowther & Smith, 1996) sits close to the threshold, as does BAT2 (WN2b(h)), for which Hainich et al. (2014) obtained a negligible hydrogen content.
At longer wavelengths, the Brackett series and adjacent He ii lines are not easily amenable to the determination of hydrogen content. However, we also consider the equivalent width ratios of He i 1.08m/He ii 1.01m (ionization) to Br/He ii 2.19m (ionization and hydrogen content) for narrow-lined WN stars in Fig. 9(c), and broad-lined WN stars in Fig. 9(d). These closely resemble the above comparisons, and we have attempted to indicate the regions of stars with atmospheric hydrogen, although high ionization, weak-lined WN3–4 stars with modest hydrogen content (e.g. WR128, WR152 Crowther et al., 1995c) are less straightforward to interpret since they lie close to the boundary. Amongst broad-lined WN stars, WR136 again lies close to the threshold obtained from narrow-line stars.
We have also utilised the He i 1.70m/He ii 1.69m ratio together with Br/2.189m He ii. In Fig. 9(e) and (f) we present a comparison between these equivalent width ratios for narrow-lined and broad-lined WN stars, respectively. Fewer stars are available for this comparison owing to the lack of HK-band spectroscopy for LMC WN stars and the difficulty in deblending He i 1.70m/He ii 1.69m, especially in broad-lined stars. Nevertheless, WN(h) and WNh stars again exhibit an excess in Br at a given ionization, once again with the notable exception of weak-lined WN3–4(h) stars. The shaded region in Fig. 9(e) indicates regions for which hydrogen is present, which are duplicated for broad-lined stars in Fig. 9(f), with WR136 (WN6b(h)) again close to this threshold.
For WN stars with exceptionally high interstellar reddenings, establishing hydrogen contents of WN stars solely from K-band spectroscopt is extremely challenging. This is illustrated in Fig. 10, in which we compare the equivalent width ratios of (N v 2.10 + He i + N iii 2.11)/He ii 2.19 to Br/He ii 2.19m. These are ratios are well correlated, although a Br excess is generally apparent in WN(h) and WNh stars (crosses and plus symbols), such that we have indicated these in the shaded region, although no attempt has been made to distinguish between modest and high hydrogen contents from these ratios.
a.2 WN/C stars
At visual wavelengths, intermediate WN/C stars usually closely resemble WN stars, albeit with prominent C iv 5808Å emission and usually weakly enhanced carbon recombination lines with respect to helium recombination lines (e.g. C iv 5471/He ii 5411). The most prominent carbon lines at near-IR wavelengths are C iv 1.19, 1.74 2.075, 2.43m, plus C iii 0.971, 2.115m. We have identified several pairs of carbon to helium lines amenable to the identification of WN/C stars at near-IR wavelengths, from spectroscopy of transition stars WR8, WR26, WR98, WR145 and BAT88.
In Fig. 11(a) we compare the ratios of C iii 0.97m/He ii 1.01m to C iv 1.19m/He ii 1.16m for WN (squares), WN/C (diamonds) and WC (triangles) stars. The intermediate carbon-to-helium abundances of WN/C stars are reflected in their high C iv 1.19m/He ii 1.16m ratios, although only late WN/C subtypes also exhibit elevated C iii 0.97m/He ii 1.01m ratios, since C iii 0.97m is not observed in early-type WN/C stars (e.g. BAT88, WN4b/CE). Nevertheless, C iv 1.19/He ii 1.16 is significantly larger in BAT88 with respect to normal WN4b stars (e.g. WR6, WR18).
The only suitable pair of carbon-to-helium recombination lines in the K-band is C iv 2.43m/He ii 2.35m, although C iv 2.07m is often also pronounced in WN/C stars (recall WR8 and WR26 in Fig. 7). We therefore consider the ratio of C iv 2.07m/(C iii + N iii + He i 2.11m), although C iii 2.115m will also be enhanced at late subtypes. Fig. 11(b) compares these ratios in WN/C stars with WN and WC stars. WN/C stars are well separated from most normal subtypes with the notable exceptions of WC9 stars which have unusually weak C iv emission lines (and strong C iii 2.11m emission), but are otherwise easily distinguished from transition WN/C stars.
Finally, C iv 1.74m/(He ii 1.69 + He i 1.70) is also a potential diagnostic, although this is complicated by the presence of Br10 + He ii 1.735 in late WN stars such that it is most effective in hydrogen-free early subtypes. By way of example, C iv 1.74m is enhanced in WR8 (WN7o/CE) with respect to WR120 (WN7o) in Fig. 7.
a.3 WC stars
For WC stars, the primary optical diagnostics involve the ratio of C iv 5808Å/C iii 5696Å, with other secondary criteria - O iii-v 5592Å for early WC stars according to Smith et al. (1990) and C ii 4267Å for late WC stars according to Crowther et al. (1998). In sharp contrast with WN stars, there are many carbon lines observed in the near-IR spectra of WC stars (Eenens et al., 1991), the strongest of which include C ii 0.990, 1.785m, C iii 0.971, 1.199, 2.115m, C iv 1.190, 1.430, 1.736, 2.075 and 2.427m. Prominent helium lines include He i 1.08, 1.70, and He ii 1.01, 1.16m, although C iv contributes significantly to lines which otherwise would be assumed to correspond to solely ionized helium, i.e. C iv (10–8) to He ii (5–4) at 1.01m.
In the Y-band, we consider a range of diagnostic ratios to discriminate between WC subtypes. Fig. 12(a) compares C iii 1.20m/C iv 1.19m to C ii 0.99m/C iii 0.97m for late (WC8–9, triangles), mid (WC6–7, diamonds) and early (WC4–5, squares) subtypes. The former cleanly separates WC8 and WC9 stars from earlier subtypes, but fails to discriminate between WC4–7 stars and is severely hindered by severe blending in broad lined systems which are common at early subtypes. Similarly, the latter separates WC9 and WC7–8 from earlier subtypes, but C ii 0.99m is extremely weak at early subtypes (absent at WC4) and lies in the electron scattering wing of C iii 0.97m for broad-lined stars.
Helium lines are not generally considered when visually classifying WC stars, since the strongest features are severely blended with carbon lines, e.g. He ii 4686Å (He i 5876Å) lies in the wing of C iii 4650Å (C iv 5808Å). However, in the Y-band, He ii 1.01m is relatively isolated – albeit with a C iv contribution - and He i 1.08m is prominent throughout the WC sequence. Therefore, Fig. 12(b) compares C iii 0.97m/He ii 1.01m to He i 1.08m/He ii 1.01m for late, mid and early subtypes. The former separates WC4 stars from later subtypes, owing to a ratio of 1, and also discriminates between WC6–7 and WC8–9 stars at a ratio of 10. The latter ratio maps reasonably well onto the WC sequence, especially for WC9 (10) and WC4 (1) subtypes.
Fig. 12(c) compares the ratio of the strongest carbon lines in the K-band, C iii 2.11m/C iv 2.07m, with the ratio of the (He ii 1.69m + He i 1.70m)/C iv 1.74m. For WC4–7 stars the latter is effectively the ratio of the helium-to-carbon recombination lines (12–7 transition for helium, 9–8 transition for carbon), so will depend upon the C/He ratio as well as ionization. In contrast, the 1.7m helium feature is dominated by He i 1.70m in WC8 and especially WC9 subtypes, so this represents a better ionization indicator for late subtypes. The K-band diagnostic cleanly separates WC4–7 subtypes from WC8 and WC9 stars, but fails to distinguish between early subtypes, as has previously been shown (Eenens et al., 1991; Figer et al., 1997). The H-band diagnostic also distingishes between WC8 and WC9 stars, but also maps earlier WC subtypes reasonably well. C ii 1.78m is also prominent at late subtypes, such that C ii 1.785m/C iv 1.736 1 in WC9 subtypes.
a.4 WO stars
At visual wavelengths, strong O vi 3818Å emission discriminates WO stars from early WC stars, with ratios of O vi to O v 5590Å and/or C vi 5808Å providing WO subclasses Crowther et al. (1998). The majority of known WO stars have been observed in the near-IR (this work, Tramper et al., 2015). Unfortunately, the sole ionization stage of oxygen present in the near-IR spectrum of WO stars is O vi (1.075m, 1.458m, 1.917m, 2.463m), preventing discrimination between WO subtypes, plus the exceptionally broad lines of WO stars leads to severe blending in most spectral regions.
Nevertheless, beyond their line widths, WO stars present a highly distinctive near-IR appearance. They can be distinguished from WC4 stars by the presence of O vi features, with high O vi 1.075m/C iv 1.190m and O vi 1.46m/C iv 1.40m ratios, although 1.46m is severely blended with C iv 1.43m and He ii 1.47m in most WO stars, and 1.075m is blended with He i 1.083m. He iv LH41-1042 Neugent et al. (2012); Tramper et al. (2015) is unambiguously characterized as a WO star from the presence of O vi features, although C iii 0.97m is prominent at this (WO4) subtype. Since O vi lines are weak, WO stars also exhibit unusually weak helium lines with respect to early WC stars, such that WO stars possess He ii 1.16/C iv 1.19 0.2, versus 0.5–0.6 in WC4 subtypes (recall Fig. 6). In addition, C iii 0.97m and C iv 2.07m are weak in WO3 stars and absent in WO2 stars.
|WR||WN||Ref||He ii 1.01||He i 1.08||P||N v 1.11||He ii 1.16||P||He ii 1.48||N v 1.55||He ii 1.69||He i 1.70||He i 2.06||N iii-v 2.11||Br||He ii 2.19||Note||Data|
|WR6||WN4b||S96||2500||2.63||— 2.60 —||2.30||2.02||2.18||0.8:||1.43||1.46||1.53||1.76||2160||2.20||2.112.10||I1,M1|
|S90: (Smith et al., 1990); C06: (Crowther et al., 2006), C11: (Crowther & Walborn, 2011), * W.D.Vacca (priv. comm.) Note: 2.10 = N v 2.100; 2.11 = He i 2.112 + N iii 2.116; 2.19 = He ii 2.189|
|WR||WN||Ref||He ii 1.01||He i 1.08||P||N v 1.11||He ii 1.16||P||He ii 1.48||N v 1.55||He ii 1.57||Note||Data|
|BAT2||WN2b(h)||F03b||2680||2.48||1.0||2.13||1.65||1.9:||He i 1.08 absent||N4|
|BAT5||WN2b||F03b||2430||2.56||1.0||1.1:||1.0:||2.20||1.66||1.9||He i 1.08 absent||N4|
|BAT17||WN4o||F03b||1780||2.29||2.22||1.23||0.8:||2.01||1.63||1.76||0.8:||1.1||log (N v 1.55/He ii 1.57) = 0||N3|
|BAT76||WN9h||C95||370||0.66||2.12||1.30||log (Br/He ii 2.189) = +1.5||C1|
|BAT88||WN4b/CE||F03b||2670||2.75||2.54||1.69||2.46||1.93||2.20||1.54||log (C iv 1.19/He ii 1.16) = –0.9||N4|
|AB1||WN3ha||F03a||1550||1.88||0.0||0.78||1.53||1.38||1.4:||He i 1.08 absent||N3,N4|
|AB4||WN6h||F03a||1070||1.94||1.45||1.20||1.61||1.61||1.41||0.9||log (Br/He ii 2.189) = –0.08||N1|
|M86: (Moffat & Seggewiss, 1986), C95: (Crowther et al., 1995a), S96: (Smith et al., 1996), W99 (Walborn et al., 1999), F03a (Foellmi et al., 2003a), F03b (Foellmi et al., 2003b), S08: (Schnurr et al., 2008) ‡We adopt WN3o for BAT47 on the owing to its modest FWHM ( 1900 km/s) whereas Foellmi et al. (2003b) adopted WN3b.|
|WR||WC||Ref||C iii 0.97||C ii 0.99||He ii 1.01||He i 1.08||He ii 1.16||C iv 1.19||C iii 1.20||C iv 1.43||He i-ii 1.70||C iv 1.74||C ii 1.78||He i 2.06||C iv 2.07||C iii 2.11||C iv 2.43||Data|