Phase-Resolved XMM-Newton and Swift Observations of Wr 25
We present an analysis of long-term X-ray and optical observations of the Wolf-Rayet binary WR 25. Using archival data from observations with the XMM-Newton and the Swift observatories spanning over yr, we show that WR 25 is a periodic variable in X-rays with a period of days. X-ray light curves in the 0.5-10.0 keV energy band show phase-locked variability, where the flux increased by a factor of from minimum to maximum, being maximum near periastron passage. The light curve in the soft energy band (0.5-2.0 keV) shows two minima indicating the presence of two eclipses. However, the light curve in the hard energy band (2.0-10.0 keV) shows only one minimum during the apastron passage. The X-ray spectra of WR 25 were explained by a two-temperature plasma model. Both the cool and the hot plasmas were constant at and keV throughout an orbital cycle, where the cooler plasma could be due to the small scale shocks in a radiation-driven outflow and the high temperature plasma could be due to the collision of winds. The column density varied with the orbital phase and was found to be maximum after the periastron passage, when the WN star is in front of the O star. The abundances of WR 25 were found to be non-solar. Optical V-band data of WR 25 also show the phase-locked variability, being at maximum near periastron passage. The results based on the present analysis indicate that WR 25 is a colliding wind binary where the presence of soft X-rays is attributed to individual components; however, hard X-rays are due to the collision of winds.
Subject headings:Star:individual (WR 25) – star:binary – star:X-ray – star:Wolf-Rayet – star:wind
Massive O-type stars evolve into Wolf-Rayet (WR) phases when their hydrogen fuel has been consumed and products of nuclear fusion appear in their atmosphere, before ending their lives as core-collapse supernovae (Doom, 1987; Smartt et al., 2009; Smartt, 2009). Typically, the progenitors of WR stars have initial masses greater than 25 (Crowther, 2007), and they spend 10% of their 5 Myr lifetime in the WR phase (Meynet & Maeder, 2005). Spectra of WR stars are predominantly characterized by emission lines of He and N (WN stars), He and C (WC stars) and He and O (WO stars). In the evolutionary sequence of massive stars, WC and WO stars are expected to correspond to later stages than WN stars. WR stars are known to produce strong stellar winds driven by their strong radiation field. The stellar winds can reach velocities up to 1000-3000 km s with mass-loss rates of yr, depending upon mass and age (Hamann et al., 2006). These winds not only affect the evolution of WR stars but also have a tremendous impact on their ambient media. The detailed physical properties of WR stars are summarized in many past reviews (Abbott & Conti, 1987; van der Hucht, 1992; Crowther, 2007).
X-ray emissions from WR binaries may consist of the combined emission from intrinsic stellar wind shocks and colliding wind shocks between the two binary components (Prilutskii & Usov, 1976; Cherepashchuk, 1976; Luo et al., 1990; Usov, 1992; Stevens et al., 1992). Colliding wind binary systems often exhibit periodic X-ray modulation either because of the change in binary separation or due to the changing circumstellar opacity along the line of sight to the collision zone, resulting from the orientation of the system with respect to the observer. Serious investigations of these variations were carried out after the advent of high-quality facilities like XMM-Newton and Chandra. Many O+O and WR+O binaries as observed with XMM-Newton, show different kinds of phase-locked modulations e.g., V444 Cyg and CD Cru (Bhatt et al., 2010a), HD 159176 (De Becker et al., 2004), HD 152248 (Sana et al., 2004), HD 93403 (Rauw et al., 2002), etc. It is believed that single WR and OB stars emit soft X-rays (kT 1 keV) via shocks that are set up by instabilities in their supersonic line-driven winds (Lucy & White, 1980; Lucy, 1982; Owocki et al., 1988). Gräfener & Hamann (2005) have shown that winds of WR stars can be driven by radiation pressure alone if multi-scattering effects are taken into account.
WR 25 (=HD 93162) is a bright (V 8.03 mag) WR binary located in the Carina Nebula region. WR 25 is classified as WN6h + O4f (van der Hucht, 2001). Gamen et al. (2006) detected periodic radial velocity variations in WR 25 and suggested that WR 25 is an eccentric binary system with a probable period of about 208 days. The basic parameters of WR 25 are summarized in Table 1. The origin of the large X-ray flux of WR 25 is suggestive of colliding wind emission in a binary system (Pollock, 1987; Raassen et al., 2003; Pollock & Corcoran, 2006). The X-ray to bolometric luminosity () ratio of (Seward & Chlebowski, 1982) for WR 25 is an order of magnitude higher than those observed for single massive stars. Raassen et al. (2003) noticed that the X-ray luminosity of WR 25 had remained relatively constant over a time span of 10 yr. Later, Pollock & Corcoran (2006) found an increase in X-ray luminosity of more than a factor of two and suggested that the observed X-ray variability is a result of colliding wind emission in a moderate eccentric binary. In this context, we use the full set of archival XMM-Newton and Swift X-ray observations of WR25 to further investigate the properties of the colliding winds in this system. Our aim also extends to search for variability in the V-band using the All Sky Automated Survey [ASAS; Pojmanski (2002)] archival data.
The paper is organized along the following lines: Section 2, describes the observations and data reduction. The light curve analysis is given in Section 3, X-ray spectral properties of WR 25 are described in Section 4. Section 5 describes the V-band optical observations. In Section 6, we present the discussion and conclusions.
|Period (d)||1||(km s)||2480||2|
|V (km s)||1||dist.(kpc)||3.24||3|
|K (km s)||1||()||10||2|
Here: - eccentricity, V - radial velocity, K - radial velocity
amplitude, - orientation of periastron, T - Julian date of
periastron passage, - semi major axis, - terminal
velocity, - mass loss rate
References: (1) Gamen et al. (2006); (2) Crowther & Dessart (1998); (3) van der Hucht (2001); (4) Seward & Chlebowski (1982)
2. Observations and data reduction
WR 25 was observed with the XMM-Newton satellite using various detector configurations on twenty occasions from the year 2000 to 2009, spanning 8.5 yr. The XMM-Newton satellite is composed of three coaligned X-ray telescopes (Jansen et al., 2001), which simultaneously observe a source, accumulating photons in three CCD-based instruments, namely the nearly-identical MOS1 and MOS2 (Turner et al., 2001) detectors and the PN (Strüder et al., 2001) detectors, which comprise the European Photon Imaging Camera (EPIC). The EPIC instrument provides imaging and spectroscopy in the energy range from 0.15 to 15 keV with an angular resolution of 4.5-6.6 arcsec and a spectral resolution () of 20-50. Exposure times for these observations were in the range of 6-60 ks. A log of observations is provided in Table 2.
The data were reduced with standard XMM-Newton Science Analysis System (SAS) software, version 12.0 using version 3.1 calibration files. The pipeline processing of raw EPIC Observation Data Files was done using the epchain and emchain tasks which allow calibrations both in energy and astrometry of the events registered in each CCD chip. We have restricted our analysis to the energy band 0.5 - 10.0 keV as the background contribution is particularly relevant at high energies where stellar sources have very little flux and are often undetectable. Event list files were extracted using the SAS task evselect. The epatplot task was used for checking for pile-up effects and no observations were affected by pile-up. Data from the three cameras were individually screened for the time intervals with high background. The observations affected by high background flaring events were excluded (see Table 2). X-ray light curves and spectra of WR 25 were generated from on-source counts obtained from circular regions with a radius of 30 around the source. The background was chosen from several source-free regions on the detectors surrounding the source. We used the tool epiclccorr to correct for good time intervals, dead time, exposure, PSF, quantum efficiency and background subtraction. The SAS task especget was used to generate the spectra, which also computes the photon redistribution as well as the ancillary matrix. Finally, the spectra were rebinned to have at least 20 counts per spectral bin.
WR 25 has also been regularly monitored by the Swift X-ray telescope (XRT) since 2007 November. The XRT observes from 0.3 to 10 keV using CCD detectors, with energy resolution of eV at keV (Burrows et al., 2005). The exposure times for the XRT observations varied from 1 to 26 ks. All XRT data were collected in photon counting mode. A log of observations is given in Table 3. We have excluded observations with exposure times less than 2 ks due to poor or no signal. In order to produce the cleaned and calibrated event files, all the data were reduced using the Swift xrtpipeline task (version 0.12.6) in which standard event grades of 0-12 were selected, and calibration files were used from the CALDB 2.8 release111http://heasarc.gsfc.nasa.gov/docs/heasarc/caldb/caldb_intro.html.
For every observation, images, light curves and spectra were obtained with the xselect (verson 2.4) package. For each observation, source spectra and light curves were extracted from a circular region with a radius of 30. For the background estimation, an annular region with inner and outer radii of 69 and 127 were used around the source region. The spectra were grouped to have a minimum of twenty counts per energy bin with grppha. The spectra were corrected for the fractional exposure loss due to bad columns on the CCD. For this, we created exposure maps with the xrtexpomap task, which is used as an input to generate the ARF with the xrtmkarf task. For the RMF, the latest version was used from the HEASARC calibration data base.
|Data||Rev.||Observation||Detector||Obs. Date||U.T.||Exposure||Bkg flare||Effective||Off axis|
M1 and M2 stand for MOS1 and MOS2, respectively.
TH and ME stand for thick and medium filters, respectively.
Exposure start time. Background proton flare duration during the
PIs of observations were Dr. Albert Brinkman, Dr. Michael Corcoran, Dr. Fred Jansen, and Dr. Kenji Hamaguchi
|Set||Observation||Obs. Date||U. T.||Exp.||offset|
|s38||00031097042||2009 06 04||06:51:00||3923.2||1.858|
Proposers of observations were F. Senziani, A. Pollock, E. Pian and M. Corcoran.
2.3. ASAS V-band observations
WR 25 was observed from 2000 December 3 to 2009 December 3 by ASAS222http://www.astrouw.edu.pl/asas. We have used only ‘A’ and ‘B’ grade data within 1 arcsec to the target of WR 25. ASAS photometry provides five sets of magnitudes corresponding to five aperture values varying in size from 2 to 6 pixels in diameter. For bright objects, Pojmanski (2002) suggested that the magnitudes corresponding to the largest aperture (i.e. the aperture diameter of 6 pixels) are useful. Therefore, we took magnitudes corresponding to the largest aperture for further analysis.
3. X-ray light curves and period analysis
The background-subtracted X-ray light curves as observed from XMM-Newton-EPIC, and Swift-XRT in the total (0.5-10.0 keV), hard (2.0-10.0 keV) and soft (0.5-2.0 keV) energy bands are shown in Figures 1 (a) and (b), respectively running from top to bottom. The light curves were binned at 2000 s intervals. The variability in each band is clearly seen. For additional confirmation of the variability in all bands, the significance of deviations from the mean count rate were measured using the - test, defined as
where is the average count rate, C is the count rate of observations and is the error corresponding to C. The statistic was compared against a critical value () for 99.9% significance level, obtained from the -probability function. For the XMM-Newton MOS light curves, values were obtained to be 6440, 4403 and 4499 in the total, hard and soft bands, respectively. These values of are very large in comparison to the of 1430 for 1599 degrees of freedom. For Swift -XRT light curves in total, hard and soft bands, values of 5188, 5113 and 5124 were obtained, respectively, with 393 for 483 degrees of freedom. This indicates that WR 25 is essentially variable in all X-ray bands.
The present long-term X-ray data as observed from XMM-Newton permit us to derive the orbital period of WR 25. We have performed a period analysis of light curves in all bands with a Lomb-Scargle periodogram (Lomb, 1976; Scargle, 1982). The plots from top to bottom in Figure2 show the power spectra in the total, hard and soft bands, respectively. The frequency corresponding to the highest peak was found to be cycles day in all bands. Other frequencies were also noticed in the power spectra but are not consistent with each other among all the energy bands. Thus, the corresponding period of days appears to be real. Also, the derived period is very similar to the orbital period (= days) derived by Gamen et al. (2006) using radial velocity measurements.
Further, the X-ray light curves as observed from XMM-Newton-MOS and Swift-XRT in the total, hard and soft energy bands were folded using the ephemeris HJD = 2451958.0 + 207.85E (Gamen et al., 2006) and are shown in Figures 3(a) and (b), respectively. The zero phase in the folded light curves corresponds to the time of periastron (HJD 2451958.0). X-ray light curves in each band from both observations show similar behavior. The light curves in the individual bands show phase-locked variability. In the soft and total-band light curves, the count rates decrease when going from orbital phase 0.0 (i.e., near the periastron passage) to orbital phase then increase up to phase . Count rates further decrease to phase before reaching a maximum value near periastron passage. However, in the hard band, the count rates were systematically decreasing when going from periastron to apastron. The ratio of maximum to minimum count rates in total, soft and hard bands were found to be 1.9, 2.6, and 1.7 for XMM-Newton-MOS observations, and 2.5, 2.6, and 2.9 for Swift-XRT observations, respectively.
The hardness ratio (HR) defined as (Hard -Soft)/(Hard+Soft) can reveal information about the spectral variations. The HR curve displayed in plot (d) of Figure 3 exhibits similar behavior to that of the light curves in the hard band. The maximum value of the HR during periastron passage indicates a harder spectrum.
4. X-ray Spectral Analysis
X-ray spectra of WR 25 as observed by the EPIC detector of XMM-Newton at different orbital phases are shown in Figure 4. Below 1 keV, the X-ray spectra were found to be affected by high extinction. Strong emission lines like Fe XVII (0.8 keV), Ne X (1.02 keV), Mg XII (1.47 keV), Si XIII (1.853 keV), S XV (2.45 keV), Ar XVII (3.12 keV),Ca XIX+XX (3.9 keV), and Fe XXV (6.67 keV) were identified in the X-ray spectra of WR 25. In order to derive the spectral parameters at different orbital phases, we performed spectral analysis of each data set corresponding to different orbital phases using simultaneous/joint fitting of EPIC data by models of the Astrophysical Plasma Emission Code [APEC; Smith et al. (2001)], as implemented in the X-ray spectral fitting package xspec (Arnaud, 1996) version 12.7.1. For spectral fitting, we adopted the similar approach of Raassen et al. (2003) and Pollock & Corcoran (2006). The form of the model used was wabs(ism)*wabs(local)*(vapec+vapec). A minimization gave the best fit model to the data. The presence of interstellar material along the line of sight and the local circumstellar material around the stars can modify the X-ray emission from massive stars. We have applied the local absorption in the line of sight to the source using photoelectric absorption cross-sections according to Morrison & McCammon (1983) and modeled as wabs with two absorption components i.e. interstellar medium () and local () hydrogen column densities. Assuming the E(B-V ) 0.63 mag (van der Hucht, 2001) and a normal interstellar reddening law towards WR 25, and using the relation given by Gorenstein (1975), the was estimated to be cm. For the first stage of spectral fitting, we fixed the and abundances of He(=2.27), C(=0.15), and N(=5.9), and varied other parameters for all phases. The values of He, C, and N abundances were adopted from the optical spectrum of WR 25 (Crowther et al., 1995). The temperatures for both components were found to be constant within the level at all phases. However, abundances of Ne, Mg, Al, Si, Ca, Ar, Fe, and Ni were found to be constant within the 1-2 level. Average values of elemental abundances and temperatures are given in Table 4. In the next stage of spectral fitting, we fixed temperatures and abundances of all elements at their average values for all phases along with , and varied , and normalizations of both components. Many spectra were taken at the same orbital phase; therefore, joint spectral fittings were performed for those spectra, which were observed within a difference of in the orbital phase. In this way, from XMM-Newton observations, we have nine data points over an orbital cycle of WR 25. The final set of the best-fitted parameters is given in Table 5.
The method used for the fitting of XMM-Newton-EPIC spectra was also applied for the fitting of spectra observed from the Swift-XRT. Since WR 25 was observed more than 40 times by the Swift-XRT and many observations were taken with short exposures, the spectra with a phase difference of 0.02 were fitted jointly. In this way, we have spectra at 12 different phases over an orbital cycle from Swift observations. For the spectral fitting of Swift-XRT data, we have fixed temperatures and abundances to values as obtained from spectral fitting of XMM-Newton-EPIC data. However, and normalizations of both components were kept as free parameters in the spectral fitting. The best fit parameters from the spectral fitting of Swift -XRT data are also given Table 5. We have compared various parameters derived from both observatories. It appears that the X-ray fluxes and other spectral parameters derived from spectral fitting of both satellite data are consistent. Comparison of spectral parameters derived from both telescopes may not be possible due to the lack of simultaneous observations. However, Plucinsky et al. (2012) has shown that the spectral parameters from both observations are in agreement within .
The X-ray fluxes of WR 25 are estimated using the cflux model in xspec and are corrected for . The X-ray luminosities of WR 25 in soft (), hard () and total () energy bands were derived by using the corresponding unabsorbed flux values and a distance of 3.24 kpc. The EMs (EMs), EM and EM, corresponding to both the cool and the hot plasma components are derived from the normalization parameters. We have plotted , , , , EM, and EM as a function of orbital phase in Figure 5 (a) and 5 (b) for observations from XMM-Newton and Swift, respectively. The maximum value of was found near periastron passage and dropped suddenly to a phase and then increased to phase . After phase , decreased to phase . It appears that peaked twice during an orbital cycle. The minimum value was observed just after periastron passage, which was times lower than the observed maximum value. A similar tendency of X-ray luminosity was also observed in the total energy band, where the maximum to minimum flux ratio was found to be . Apart from the periastron passage, no other peak of was noticed in an orbital cycle of WR 25. The minimum value of was observed near phase and was times lower than that observed during the periastron passage. The value of was observed to be maximum after the periastron passage at phase 0.03. However, the minimum value of was observed at phase . The EMs, EM and EM corresponding to the cool and the hot temperature components were also found to be phase dependent, with maximum at periastron passage and minimum at apastron passage.
5. Optical V-band light curve
The V-band light curve of WR 25 as obtained from the ASAS archive is plotted in Figure 6 (a). In ASAS observations, error bars are very large (i.e. 0.035 to 0.1 mag). Therefore, it is difficult to search for any small scale variability that might be present. In order to reduce the short term fluctuations, a moving average of 10 data points in the forward direction was performed. The light curve of the moving-averaged data points is shown in Figure 6 (b). The V-band light curve does not appear to be constant over the time span of the observations; however, variations are not statistically significant. The was found to be 270 for 956 degrees of freedom (see Equation 1). This value of is very low in comparison to the of 826 corresponding to the 99.9% significance level. Furthermore, we have folded the V-band data with phase bins of 0.1 and using the ephemeris given by Gamen et al. (2006). The folded V-band light curve of WR 25 is shown in Fig 6 (c). It appears that variability is present in the light curve. The variability amplitude was found to be mmag. The maximum brightness was seen near periastron passage and after that the brightness decreased toward apastron passage, being minimum near phase 0.75.
6. Discussion and conclusions
We have carried out X-ray and optical studies of WR 25 using X-ray data from the XMM-Newton and Swiftsatellites and optical data from the ASAS archive. The X-ray spectra of WR 25 at all orbital phases are well-explained by a two-temperature plasma model. The temperatures of both components were found to be constant throughout an orbital cycle. The temperature values are similar to those derived by Pollock & Corcoran (2006) and Raassen et al. (2003) using XMM-Newton data. The temperature of the cool component (=0.628 keV) of WR 25 is comparable to that of other WN-subtype WR+O binaries: for example WR1 [0.56-0.67 keV; (Ignace et al., 2003)], WR 22 [0.6 keV; (Gosset et al., 2009)] WR 47 [ keV; (Bhatt et al., 2010a)], WR 139 [ (Bhatt et al., 2010a)], and WR 147 [0.7-0.8 keV; (Skinner et al., 2007)]. A plasma temperature of 0.6-0.8 keV is also dominant for massive OB stars (Sana et al., 2006; Nazé, 2009; Bhatt et al., 2010b). The soft X-ray component may originate in the winds of the individual components of WR 25 via radiation driven instability shocks. The ratio between the wind momentum () and the momentum of the radiation field () for WR 25 was derived to be (Hamann et al., 2006), indicating that wind of WR 25 is driven by radiation pressure. Furthermore, the derived temperature of the cool component appears to be realistic for radiation-driven wind shocks. The derived plasma temperature for WR 25 can be used to estimate the pre-shock velocity using the relation keV (Luo et al., 1990), where is mean mass per particle in units of the proton’s mass (1.16 for a WN star and 0.62 for an O-type star), and is shock velocity in units of 1000 km s. The pre-shock velocities corresponding to the cool temperature for the WN and O components of WR 25 were found to be 527 and 721 km s, respectively. These values are about a factor of more than those predicted by the radiative shock model of Lucy (1982). The advanced version of the wind-shock model by Owocki et al. (1988) predicts X-ray emission up to 1 keV. The hydrodynamic shocks, which are expected to occur within unstable stellar winds of massive O-type stars, also show similar pre-shock velocities (Feldmeier et al., 1997). Observationally some apparently single WN stars (e.g. WR 1, WR 6, and WR 110) show intrinsic X-ray emission (Skinner et al., 2002a, b; Ignace et al., 2003) while apparently single WC stars and the WN-subtype WR star WR 40 have not been detected in X-rays thus far (Oskinova, 2005; Gosset et al., 2005).
The temperature corresponding to hot the plasma of WR 25 is intermediate compared to those observed for other similar WR binaries (Zhekov, 2012). Using the mean particle weight for WN stars, the maximum shock temperature on the line of centers for an adiabatic shock corresponds to a pre-shock wind velocity of km s (see Luo et al. 1990), which is % of the observed terminal velocity (Crowther & Dessart, 1998; Niedzielski & Skorzynski, 2002) of WR 25. This could be due to an oblique wind collision over most parts of the shock, which occurs with a lower velocity normal to the shock and thus leads to a lower plasma temperature. The derived values of and X-ray luminosities exhibit phase-related time variability. We found to be maximum after periastron passage and minimum when is maximum during periastron passage. It also appears that reached a peak value during the eclipse, indicating that the extra absorption could be due to winds from the WN star. The column density along the line of sight through the WR wind to the region of stellar wind collision is (Usov, 1992), where is the distance between the binary components. Using the parameters of WR 25, was estimated to be cm, which is quite similar to the observed maximum value (see Table 5).
|s28||0.163||0.52||8.22||3.22||6.13||3.27||9.41||1.17 ( 49)|
|s37,s38,s39,||0.306||0.32||6.34||2.96||7.76||3.01||10.77||1.12 ( 85)|
|s29||0.476||0.23||4.90||2.32||7.47||2.38||9.85||1.16 ( 31)|
|s01||0.703||0.46||7.18||3.13||6.05||3.07||9.12||1.50 ( 20)|
|s02||0.889||0.40||11.86||4.59||11.48||4.76||16.23||1.11 ( 25)|
|s03,s04,s05,||0.924||0.40||13.19||5.89||13.14||5.95||19.10||1.05 ( 93)|
|s11,s12,s13||0.975||0.56||11.25||7.74||9.04||7.22||16.26||1.14 ( 92)|
|s14, s31||0.995||0.77||10.73||7.11||5.87||6.47||12.34||1.15 ( 88)|
Notes: EM and EM are in units of cm, is in units of cm, and , and are unabsorbed X-ray luminosities in soft, hard and total bands in units of 10 . Abundances are in units of solar photospheric values. is per degree of freedom and dof is the degree of freedom.
The EMs corresponding to the cool and hot temperatures change substantially, reflecting the variations in X-ray luminosities in the soft and hard bands. The X-ray luminosities observed for WR 25 were found to be more than that for other close WR+O binaries and single WN stars (Zhekov, 2012; Skinner et al., 2010). X-ray light curves of WR 25 as observed from XMM-Newton and Swift show similar behavior. However, in terms of phase coverage, the light curves from Swift are much better than those observed from XMM-Newton. The deficit in X-ray flux just after periastron passage could be due to the eclipse of the wind interaction zone by the wind of the WN star. In all bands, the excess emission is strongest near periastron passage. The X-ray enhancement after the eclipse in the soft band indicates that besides the colliding wind, individual components of the WR 25 system also contribute to the X-ray emission. The soft X-ray flux peaked near phase 0.5, which further supports that both of the components of WR 25 are sources of soft X-rays, enhancing the combined soft X-ray flux when both of the stars are completely visible to the observer. During phase , the soft X-rays further decrease to the minimum value indicating the possibility of a secondary eclipse, when the O-star is in front of the primary WN star. The deeper primary eclipse at phase could be due to the larger opacity of the WN wind. On the other hand, was found to be at minimum during the phase . Both components of WR 25 are farthest apart at phase 0.5; therefore, the collision is weak, generating fewer X-rays. This indicates that the hard X-rays originate from the collision zone of the wind giving enhanced flux during periastron passage. Stevens et al. (1992) and many other authors showed that strong winds from massive stars collide and generate hard X-rays, in addition to softer X-ray components due to intrinsic or individual components (Berghoefer et al., 1997; Sana et al., 2006). The phase-locked X-ray variability could be a result of changing separation in an eccentric orbit of WR 25. Fig 7 shows that the X-ray luminosity varies as a function of the inverse binary separation [1/(D/a)]. It appears that soft X-ray luminosity does not depend on the binary separation. However, it is clearly seen that the hard X-ray luminosity is linearly dependent on the inverse of the binary separation. The Pearson correlation coefficients for versus 1/(D/a) and versus 1/(D/a) were derived to be -0.02 and 0.87 with a probability of no correlation found to be 0.94 and , respectively. This indicates that the hard X-rays in WR 25 are due to the collision of winds. The wide massive binary systems whose X-ray luminosities follow the 1/D relation are Cyg OB2 # 9 (Nazé et al., 2012) and HD 93205 (Antokhin et al., 2003). Cyg OB2 # 8A and WR140 are the massive binary systems which deviate from expectations at periastron, since the collisions then become radiative (De Becker et al., 2006; Pollock, 2012; Corcoran et al., 2011). However, there are a few wide binary systems that do not follow the 1/D variation e.g., WR 11 (Rauw et al., 2000), WR 22 (Gosset et al., 2009). The present X-ray light curves of WR 25 strongly support the idea that during the rise in X-ray emission around periastron passage the X-ray emission is primarily due to colliding winds (Willis et al., 1995; Stevens & Pittard, 1999).
Using V mag, a distance of kpc, a bolometric correction of -4.5 (Hamann et al., 2006) a,nd an anomalous reddening of 4.5 mag (Hamann et al., 2006), the bolometric luminosity of WR 25 is calculated to be . The maximum and minimum values of for WR 25 are thus calculated to be and , respectively. The derived value of during periastron passage is similar to that derived by Seward & Chlebowski (1982) using Einstein observations. During the phases and , the possible primary and secondary eclipse positions where the WN and O-type stars contributed more in soft X-rays, respectively, the are calculated to be and , respectively. These values of are similar to those obtained for other single WR and O-type stars (Skinner et al., 2010; Sana et al., 2006).
The hard X-rays from WR 25 appear to be a result of collision of winds from binary companions. The gas in the colliding wind regions could either be adiabatic or radiative depending on the cooling parameter () as (Stevens et al., 1992), where is in units of 10 km s, is the distance from the contact to star in units of km and is in units of yr. For , the wind can be assumed to be adiabatic while for , it is roughly isothermal. The cooling parameter is directly proportional to the distance of contact (); as a consequence of this, (Stevens et al., 1992). This means that the shocked region will be adiabatic for longer period binaries. WR 25 is one of the longer orbital period binaries, indicating the shocked region is adiabatic. Furthermore, if we assume that the contact of winds is close to the O-type star of WR 25, the value of is calculated to be , indicating an adiabatic wind in the shocked region. In our analysis, the abundances of O, Ne, Mg, Si, S, Ca, Ar, Fe and Ni were found to be non-solar, which is expected from WR stars because they are in advanced evolutionary stages. The non-solar abundances of these elements could be due to the presence of hot plasma near the surface of the O4 star, due to the collision of strong supersonic winds from the WN6 star with the less powerful supersonic wind from O4 star. For wide binaries, Pittard & Stevens (2002) showed that WR winds dominate the X-ray luminosity. Luo et al. (1990) and Myasnikov & Zhekov (1993) have also found that the shocked WR stellar wind dominates the X-ray emission for WR+O systems in the adiabatic limit. Furthermore, the shocked volume and the emission measure near the O-type star may be dominated by WN winds; therefore, the X-ray emitting plasma also shows non-solar abundances. When we compared the derived values of abundances of different elements for WR 25 with other WN-type stars, we found that the abundances of Ne and S in WR 25 are similar to that of WN-type stars (Smith & Houck, 2005; Ignace et al., 2007)
The light curve of WR 25 in the V band shows minimum light near phase 0.75, which could be possible position of eclipse when the WN star is behind the O-star. It appears that the WN star is brighter than the O-type star in the WR 25 system; therefore, may observe an eclipse near phase 0.75. Hamann et al. (2006) also showed that O-type stars contribute a minimum of 15% to the total light of the system in the V band and Gosset et al. (1994) had also noticed the photometric variation in WR 25 with a variability amplitude of 0.02 mag in the Strömgren b-band on a time scale of 250 days. Many other WR binaries also show phase-locked variability in the optical band and the variability is attributed to their binary nature (Lamontagne et al., 1996). The presence of eclipses in X-ray and optical light curves indicates that the orbital plane of WR 25 has a high inclination angle. The minimum of the V-band light curve is consistent with the phase minimum at of the soft X-ray light curve. However, unlike to the soft X-ray light curve we could not see any other minimum in the V-band light curve near phase 0.03, where was also found to be maximum. This could be due to the poor phase coverage of the V-band light curve.
The analysis of the present data shows that WR 25 is a colliding wind binary with an orbital period of days, where the hard X-rays could be due to the collision zone while soft X-rays could be attributed to individual components.
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