High Energy Processes in Wolf-Rayet Stars

TL;DR

This paper reviews high-energy processes in Wolf-Rayet stars, emphasizing X-ray emissions, and presents new Chandra observations of the WR+O binary CQ Cep to test emission models.

Contribution

It provides a summary of high-energy phenomena in WR stars and offers new observational data to evaluate X-ray emission models.

Findings

Chandra observations of CQ Cep cover a full orbit.

Data tests existing X-ray emission models.

Enhanced understanding of WR star winds and emissions.

Abstract

Wolf-Rayet (WR) stars are massive ($\geq$10 M$_{\odot}$) evolved stars undergoing advanced nuclear burning in their cores, rapidly approaching the end of their lives as supernovae. Their powerful winds enrich the interstellar medium with heavy elements, providing raw material for future generations of stars. We briefly summarize high-energy processes in WR stars, focusing mainly on their X-ray emission. We present new results from Chandra observations of the eclipsing WR+O binary CQ Cep covering a full orbit which stringently test X-ray emission models.

Tables

Table 1. Table 1: Chandra Observations of CQ Cep

ObsId	Start Date	Duration	Phase	Rate	Pvar
		(ks)	($\varphi$)	(c ks-1)
14538	2013 Mar 19	85.6	($-$0.10, 0.50)	24.7$\pm$5.6\tnote$†$
17734	2017 Feb 27	18.3	($-$0.07, 0.06)	20.9$\pm$4.0	0.16
20016	2017 Mar 4	16.4	(0.66, 0.78)	20.3$\pm$3.9	0.07
20017	2017 Mar 5	33.9	(0.43, 0.67)	20.5$\pm$4.1	0.05
20018	2017 Mar 5	19.2	(0.79, 0.93)	20.6$\pm$4.7	0.04

Table 2. Table 2: CQ Cep Spectral Fits (shock model)

Year	NH,1, NH,2	kT1, kT2	Fx\tnote$†$	log Lx\tnote$‡$
	(1021 cm-2)	(keV)		(erg s-1)
2013	4.4[1.7], 0.3[0.1]	0.6[0.1], 3.4[0.6]	2.09	33.43
2017	6.5[2.7], 0.3[0.1]	1.0[0.3], 3.6[0.9]	2.20	33.53

Full text

††articletype: Article Type

1 Wolf-Rayet Stars: Overview

WR stars are broadly classified into three subtypes based on their optical spectra: nitrogen-type (WN), carbon-type (WC), and oxygen-type (WO). WO stars are the most evolved and only a few are known in the Galaxy. WR stars have high effective temperatures $>$20,000 K and tremendous winds with typical mass-loss rates $\dot{M}$ $\sim$10*-5* M*⊙* yr*-1*. Terminal wind speeds for WN and WC stars are typically V*∞* $\sim$1000 - 2500 km s*-1*, and even higher for WO stars. Most, but not all, WR stars are strong X-ray sources. Their X-ray emission is usually attributed to shocks associated with their supersonic winds, but shock model predictions are in some cases not compatible with the observed X-ray properties. Nonthermal (synchrotron) radio continuum emission has been detected in some WR binaries. WR binaries are potential sources of $\gamma$-ray emission and sites for Galactic cosmic ray acceleration.

2 X-Rays

2.1 Single WR Stars

X-rays have been detected from putatively single (non-binary) WN and WO stars, but surprisingly not from WC stars. X-ray luminosities of WN stars span a range of about two orders of magnitude with typical values log Lx $\sim$ 1032±1 ergs s*-1* (Skinner et al. 2012). The only galactic WO star detected so far in X-rays is WR 142 with log Lx $\approx$ 31.3 ers s*-1* at its GAIA DR2 distance of 1.74 kpc (Sokal et al. 2010; Skinner et al. 2018). WC stars are X-ray faint or X-ray quiet for reasons not yet known, but strong wind absorption may be partially responsible. The most stringent upper limit obtained for a single WC star so far is from a Chandra observation of WR 135 (WC8) which gives log Lx $\leq$ 29.99 ergs s*-1* at its GAIA DR2 distance of 2.11 kpc (Skinner et al. 2006).

The X-ray spectra of single WN stars can be acceptably modeled as a two-temperature (2T) optically-thin plasma with a cool component at kT1 $\approx$ 0.3 - 0.7 keV [T1 $\approx$ 3 - 8 MK] and a hotter component at kT2 $\approx$ 2 - 5 keV [T2 $\approx$ 20 - 60 MK] (Skinner et al. 2010; 2012). Such 2T models do not fully reflect the true physical conditions since the X-ray plasma is distributed over a range of temperatures.

The temperature of the cool component is consistent with predictions of radiative wind shock models, which posit X-ray production from shocks that form in the wind as a result of line-driven instabilities (e.g. Lucy & White 1980). Such models have had success in explaining the soft X-ray emission of some O-type stars, but their relevance to WR stars with much higher mass-loss rates and wind speeds remains to be determined. But the hotter plasma, which is prominent in single WN star spectra and in the spectrum of the WO star WR 142, is not anticipated from radiative shock models. Its origin is so far unexplained.

2.2 WR Binaries

Massive WR binaries are typically strong X-ray sources. This includes WC systems such as $\gamma^{2}$ Vel (= WR 11; WC8$+$O7) and WN systems such as WR 147 (WN8$+$B0.5V). In some cases, the X-ray emission is bright enough to obtain high-resolution X-ray grating spectra, allowing individual emission lines to be identified and studied (e.g. $\gamma^{2}$ Vel, Skinner et al. 2001; WR 140, Pollock et al. 2005; WR48a, Zhekov et al. 2014). Line information, when available, constrains the plasma temperature distribution, metal abundances, and distance from the star(s) where the line forms.

The X-ray emission of WR binaries is potentially an admixture of multiple components including that of the individual stars (and their winds) and colliding wind (CW) shock emission originating between the stars, or near the surface of the star with the weaker wind (Usov 1992). In most cases, these different components cannot be spatially-resolved with current generation X-ray telescopes.

The maximum  CW shock temperature for an adiabatic shock is kTcw $\approx$ 1.96$\mu$[V*⟂/1000 km/s]2 keV. Here, $\mu$ is the mean atomic weight (amu) in the wind ($\mu$ $\approx$ 4/3 for He-dominated WN winds) and V⟂* is the wind velocity component perpendicular to the shock front. The hottest plasma is predicted to lie on or near the line-of-centers where V*⟂* $\approx$ V*∞, if the winds have reached terminal speeds. This latter condition will be satisfied in wide binaries (Porb* $\sim$ years) but not in close binaries (Porb $\sim$ days). For typical WR wind speeds V*∞* $\approx$ 1000 - 2500 km/s, maximum shock temperatures kTcw,max $\approx$ 2 - 12 keV are expected. Temperatures in this range are indeed observed in some WR binaries, but also in some (apparently) single WR stars.

3 The Eclipsing WR Binary CQ Cep

CQ Cep (= WR 155) is an eclipsing WN6$+$O9 binary system in a near-circular high inclination 1.64 d orbit (Demircan et al. 1997). The masses and radii of the two stars are nearly equal (M*∗* $\approx$ 21 M*⊙, R∗* $\approx$ 8 R*⊙). Their separation is $D$ $\approx$ 20 R⊙*, placing the two stars nearly in contact (Demircan et al. 1997). At such close separation, the winds will not have reached terminal speeds before colliding. The higher momentum of the WR wind will overpower the O star wind and the CW shock will form at or near the O star surface. CQ Cep is a superb system for testing CW model predictions at close spacing where the winds will be at subterminal speeds and radiative cooling may be important (Stevens et al. 1992).

We have observed CQ Cep with the Chandra X-ray Observatory using the ACIS-S CCD imaging spectrometer over a full orbit (Table 1). The first half of the orbit was observed in a single uninterrupted observation in March 2013, capturing the O star passing in front of the WR star at $\phi$ = 0 (Skinner et al. 2015). The second half was observed in four exposures during Feb.-March 2017, again capturing the O star in front as well as the WR star in front. Simultaneous optical light curves were obtained using the Chandra Aspect Camera Assembly (ACA). The main objective was to search for X-ray variability during eclipses, which is expected if the hottest X-ray plasma is confined to the region on or near the line-of-centers between the two stars.

Both the primary and secondary visual eclipses are clearly seen in the optical light curves (Fig. 1-top panels). Comparison of the overlapping ACA light curves indicates that the system was slightly brighter in 2013 than in 2017. Optical variability is also apparent from significant scatter in the high-precision GAIA photometry at similar phases but different epochs. The times of minimum ACA optical brightness at $\phi$=0 in 2013 and 2017 give an accurate orbital period P = 1.641239 ($\pm$8.0e$-$07) d. Analysis of historical data suggests that the period may be variable (Koenigsberger et al. 2017).