Received 2001 October 29; accepted 2002 April 4

Transcription

1 The Astrophysical Journal, 574: , 2002 August 1 # The American Astronomical Society. All rights reserved. Printed in U.S.A. BROADENING OF NEARLY NEUTRAL IRON EMISSION LINE OF GX OBSERVED WITH ASCA Takao Endo, 1,2 Manabu Ishida, 3 Kuniaki Masai, 3 Hideyo Kunieda, 1 Hajime Inoue, 1 and Fumiaki Nagase 1 Received 2001 October 29; accepted 2002 April 4 ABSTRACT We present results from detailed analyses of the ASCA data of the X-ray pulsar GX during apastron, periastron, and intermediate orbital phases. We find that the iron K emission line possesses a nonzero energy width of ev in Gaussian during all of the three orbital phases. The observed range in the width is far larger than that expected from velocity shifts in the stellar wind. Quantitative evaluation of the width indicates that the iron K line of GX is emitted from a region within cm from the neutron star via fluorescence. In addition, we have also discovered a broad wing with 135 ev in the profile of the iron K line during the periastron phase. Since the wing profile is modulated according to the spin phase of the pulsar, we conclude that the wing is composed of blue- and redshifted 6.4 kev iron lines emanating from the bipolar accretion columns. We find that, within the accretion column, the He ii ionization front is formed at 10 7 cm from the pulsar, outside of which iron can be kept in a nearly neutral ionization state. The observed equivalent width of the broad wing component (100 ev) can be explained by emission in the accretion columns outside the He ii ionization front. Subject headings: line: profiles pulsars: individual (GX 301 2) stars: neutron X-rays: binaries X-rays: individual (GX 301 2) 1. INTRODUCTION The hard X-ray source GX was discovered as a moderately intense source by a balloon observation (Ricker et al. 1973). Soon after, coherent pulsations from the source were discovered by the SAS 3 satellite at a period of 700 s (White et al. 1976). Since then, the nature of GX has been investigated at various wave bands. The mass-donating secondary star, WRA 977, was first examined in detail by Parkes et al. (1980). They classified WRA 977 as a B2 Iae spectral type, indicating that the star has already evolved off the main sequence and is at least 10 6 yr old. This implies that the GX 301 2/WRA 977 system belongs to the class of high-mass X-ray binaries (HMXBs). Spectroscopic observations imply that the distance to WRA 977 is 1:8 0:4 kpc, although a different estimate of 5.3 kpc was proposed later on (Kaper et al. 1995). The mass-loss rate is estimated to be in the range to M yr 1 (Parkes et al. 1980; Kaper et al. 1995). The stellar wind velocity of WRA 977 measured from the hydrogen Balmer lines is 300 km s 1 at a distance of r=r c 3 (Parkes et al. 1980), where R c is the radius of the secondary, which suggests a terminal wind velocity v 1 of 400 km s 1. The large mass-loss rate and the relatively slow wind velocity result in stagnation of large amounts of matter in the binary system surrounding the neutron star. The orbital period of 41 days was established from the time of arrival (TOA) of the pulse (White, Mason, & Sanford 1978) using data obtained with Ariel 5 over a 4 yr period. The orbital parameters were then refined by Sato et al. (1986) by combining the TOA data obtained by Ariel 5, SAS 3, and Hakucho. They estimated the X-ray mass function to be 31.9 M (Sato et al. 1986; Koh et al. 1997) and a minimum companion mass to be 30.3 M, assuming a neutron star mass of 1.4 M. The lack of an X-ray eclipse combined with R c ¼ 43 R (Parkes et al. 1980) constrains the inclination to be less than 78. The binary system possesses the longest orbital period (41 days) and the highest eccentricity (e ¼ 0:46) among the persistent HMXB pulsars (excluding Be binaries). The high eccentricity results in large changes in the mass accretion rate onto the neutron star with orbital phase, producing periodic outbursts, which were first discovered by Watson, Warwick, & Corbet (1982). The X-ray luminosity varies in the range ð2 200Þ10 35 ergs s 1. It was found, however, that the peak of the outbursts appears at an orbital phase 1.4 days before the periastron passage of the neutron star (Sato et al. 1986). Several models have been proposed to explain these outbursts (e.g., Haberl 1991). In addition, BATSE observations (Pravdo et al. 1995; Koh et al. 1997) revealed the presence of secondary outbursts near the apastron passage. The X-ray intensity modulation with the binary revolution is very complicated and has not been fully understood. From a long-term pulse period history of GX 301 2, the neutron star is known to be randomly spinning up and down on various timescales, from hours to years (Nagase 1989; Bildsten et al. 1997), indicating that GX belongs to the so-called wind-fed X-ray pulsar. In such pulsars, mass accretion onto the neutron star takes place through direct capture of the stellar wind of the OB supergiant companion star. According to the classical estimate, material in the wind with impact parameter less than R acc ¼ 2GM v 2 ð1þ rel 1 The Institute of Space and Astronautical Science, Yoshinodai, Sagamihara, Kanagawa , Japan. 2 Mitsubishi Electric Corporation, Information Technology R&D Center, Ofuna, Kamakura, Kanagawa , Japan. 3 Tokyo Metropolitan University, 1-1 Minamiosawa, Hachioji , Japan; ishida@phys.metro-u.ac.jp. 879 is accreted onto the neutron star (Hoyle & Lyttleton 1939; Bondi & Hoyle 1944; Davidson & Ostriker 1973), where M is the mass of the neutron star and v rel is the relative velocity of the stellar wind with respect to the neutron star motion. Blondin et al. (1990) and Blondin, Stevens, & Kallman

2 880 ENDO ET AL. (1991) carried out detailed numerical simulations of the stellar wind behavior near the accreting compact object by accounting for gravitational and tidal forces, radiation pressure of the companion star, and heating by the pulsar radiation. Although the resultant accretion geometry is rather complicated compared to the classical picture as a result of, for example, the formation of an accretion wake trailing the neutron star, they found that the mass accretion rate agrees with that predicted by equation (1) to within a factor of 2. A number of models have been proposed to explain the X-ray spectra of HMXB pulsars, including the conventional power-law model with a high-energy exponential cutoff (White, Swank, & Holt 1983) and those that account for Comptonization in the postshock accretion (Lamb & Sanford 1979; Sunyaev & Titarchuk 1980; Becker & Begelman 1986). Kretschmar et al. (1997) applied these models to the spectrum of Vela X-1 obtained by HEXE and TTM on Mir- Kvant and concluded that the power law with an e-cut model generally provides the best representation to the observed spectra. The power-law model with an e-cut has been applied also for GX (White et al. 1983; Makishima & Mihara 1992; Mihara 1995; Orlandini et al. 2000). All of their results are consistent with a power-law photon index of ¼ 1:1 1:5 and cutoff and e-folding energies of E c ¼ kev and E f ¼ 6 12 kev, respectively. As the extension of the Comptonization models, on the other hand, Mihara, Makishima, & Nagase (1998) proposed a new model consisting of an intrinsic power law (negative photon index) plus its Comptonized component that is represented by /E expð E=kTÞ with 2 (NPEX model) with a cyclotron resonance absorption feature (see also Mihara 1995; Makishima et al. 1999). They claimed that this model provides a better representation for the spectra of HMXB pulsars obtained with Ginga compared to the power-law with e-cut model. By comparing the parameters of these two models, they showed that E c has a tight correlation with the cyclotron absorption energy E a, indicating that the conventional exponential cutoff is attributed to the cyclotron resonance absorption structure. The cyclotron absorption energy of GX is directly measured to be at E a ¼ 37:0 0:9 kev from a Ginga observation (Makishima & Mihara 1992; Mihara 1995), which is in line with that expected from the correlation. This E a results in the surface magnetic field of the neutron star of ð3:2 0:1Þ10 12 G. A remarkable characteristic of the GX spectrum is heavy and variable photoelectric absorption and a strong iron K emission line (White & Swank 1984). From Tenma observations, the time-averaged iron K emission line and iron K edge energies were found to be 6:46 0:05 and 7:36 0:05 kev, respectively. The equivalent width of the iron line varies from 200 to 1800 ev, while the hydrogen column density N H ranges from to as high as 1: cm 2 with orbital phase (Leahy et al. 1988, 1989b; Leahy & Matsuoka 1990). The intensity of the iron K line and N H are well correlated and can be interpreted as the iron K line being emitted via fluorescence from matter of solar iron composition surrounding the neutron star with 4 solid angle (White & Swank 1984; Leahy et al. 1988; Haberl 1991). However, the line central energy and the edge energy indicate a somewhat different ionization state of iron, which lead Leahy et al. (1989b) to conclude that the gas along the line of sight is cooler than the bulk of the X-ray irradiated gas. The fluorescence picture of iron K emission line is further supported by Ginga observations. Pulsations were not detected from the time series data of the iron K photons (Tashiro et al. 1991), although it is generally seen in the energy band above 2 kev. Moreover, from the power spectral analysis, Tashiro et al. (1991) discovered that the distance of the fluorescing matter should not greatly exceed several light-seconds or cm from the neutron star. In this paper we analyze all of the existing ASCA data of GX and present the results with primary emphasis on the iron K emission line. In x 2 we describe how the observations were conducted. In x 3 analysis procedures and results are summarized. The implications are discussed in x 4, and we conclude in x 5. Part of the data presented in this paper have already been published by Saraswat et al. (1996). They analyzed the data taken at around the orbital phase of 0.31 (see x 2) and found that the energy spectrum above 2 kev requires the partial-covering absorber model. 2. OBSERVATION 2.1. ASCA Observation of GX ASCA (Tanaka, Inoue, & Holt 1994) is equipped with four X-ray telescope (XRT; Serlemitsos et al. 1995) modules, with two types of X-ray detectors mounted at their focal planes: two Gas Imaging Spectrometers (GISs; Ohashi et al. 1996; Makishima et al. 1996) and two Solid-State Imaging Spectrometers (SISs; Burke et al. 1994). The ASCA observations of GX were performed in 1994 and 1996 at three different orbital phases. The observation log is summarized in Table 1. A schematic drawing of the binary orbit, with the observation dates and orbital phases, is shown in Figure 1. The first observation was performed on 1994 February 13 at an orbital phase around 0.31 with a net exposure time of 40 ks. The following two observations were carried out within the same orbit of the binary on 1996 January 26 ( 0:47) and February 16 ( 0:97) with exposure times of 20 ks each. Note that ¼ 0 is defined to be the periastron passage of the neutron star around WRA 977. Hence, we hereafter designate these three observations as intermediate, apastron, and periastron observations. We retrieved these ASCA data from the archival database at the NASA Goddard Space Flight Center (GSFC). Throughout the three observations, the GIS was operated in PH mode with the standard telemetry bit assignment, with time resolutions of 62.5 ms and 0.5 s for the high and medium telemetry editing rates, respectively. For the SIS, both the faint and the bright modes were used. Prior to data screening, we reformatted all of the faint mode data into bright mode for unified data handling. The time resolution of the SIS data is 4 s, irrespective of the telemetry editing rates. These time resolutions are much shorter compared to the pulse period (676 s) of GX Throughout the three observations, the SIS was operated in 1-CCD mode Data Selection Criteria Prior to all analyses described below, we screened the data according to the following selection criteria. We accumulated photons detected while the cosmic-ray cutoff momentum is greater than 6 GeV c 1. Data taken while the elevation angle of the pointing direction from the Earth rim was less than 5 for GIS and less than 10 for the SIS were

3 TABLE 1 Journal of the ASCA Observations Observation Observation Time (UT) Observation Mode Orbital Phase a Spectrometer High Medium Exposure b (ks) Count Rate c (count s 1 ) Period P spin (s) Midtime (MJED) Intermediate Feb 13 11: : SIS Bright Bright , GIS PH nominal PH nominal Apastron Jan 26 09: : SIS Faint Bright 13.9/ , GIS PH nominal PH nominal Periastron Feb 16 09: : SIS Faint Bright 5.1/ , GIS PH nominal PH nominal a ¼ 0 is the periastron passage of the neutron star. b Net exposure after the data screening. c Photons are accumulated with circular regions centered on GX with radii of 4 0 and 6 0 for the SIS and the GIS, respectively.

4 882 ENDO ET AL. Vol. 574 Fig. 1. ASCA observation epochs drawn on the orbit of the neutron star. The orbital parameters are taken from Koh et al. (1997). not used. In addition, we discarded the SIS data taken while the elevation angle from the bright Earth rim was less than 20. We required that the pointing direction was within less than 0=01 from the target location. We further required the Radiation Belt Monitor count rate to less than 200 counts s 1. The net exposure times listed in Table 1 are those after the data screening Soft X-Ray Sources near GX GX is located nearly on the Galactic plane and surrounded by several other sources as clearly seen in Pravdo et al. (1995). We list these sources in Table 2 in the order of increasing right ascension. Among the surrounding sources, source 1 is separated by only 3<2 from GX and is resolved only by the SIS. Sources 2 and 3 are separated from GX roughly by 8 0. Keeping these facts in mind, we accumulated photons from GX in circular regions with radii of 4 0 and 6 0 for the SIS and the GIS, respectively. Although these integration regions are essentially free of photons from source 2 through 4, they include those from source 1. To estimate the amount of contamination from source 1, we carried out energy-resolved image analysis. Fortunately, the X-ray spectrum of source 1 is found to be very soft, whereas that of GX is hard (Fig. 4). The intensity of source 1 is only 6% of GX 301 2in the band kev and is much smaller at higher energies. Therefore, the contamination due to source 1 can be neglected as long as we limit our analysis to the band above 2 kev. For background regions, on the other hand, we adopted the residual area of the same CCD chip employed for the GX observations that is out of the source integration region for the SIS and the annulus centered on GX with the inner and outer radii of 11 0 and 13 0, respectively, for the GIS. After correcting for the areal difference between the source and the background regions, we subtracted the background spectra from the source spectra for both the SIS and the GIS. Note, however, that a significant number of photons from GX fall into the background region for the SIS. We corrected for this photon leakage in the response function. We first calculated the effective areas of the XRT for both the source and the background regions. After correcting for the difference between the integrated detector areas as in the background subtraction, we then subtracted the effective area of the background region from that of the source region. Finally, the resultant effective area is reflected in the response matrix of the SIS, which we use in the subsequent spectral analyses. 3. ANALYSIS AND RESULTS 3.1. Temporal Analysis According to the selection criteria described above, we have made background-subtracted light curves for all three observations separately. The average X-ray intensities are found to be 0.96/1.06, 3.15/3.18, and 4.51/6.37 counts s 1 per SIS/GIS detector for the intermediate, apastron, and periastron observations, respectively (Table 1). We first obtain the average pulse period during each observation. Since all three ASCA observations were rather short, covering only 0.02 of the orbital cycle (Table 1), we could not determine the orbital parameters by these data alone. However, the orbital ephemeris of GX has been well established (Sato et al. 1986; Koh et al. 1997). Adopting their values, we are able to correct for the TOA of each photon to the center of mass of the binary. We then obtain the barycentric pulse periods by the epoch-folding analysis, which we find to be P spin ¼ 675:76 0:07, 677:35 0:10, and 678:08 0:19 s for the intermediate, apastron, and periastron observations, respectively (Table 1). The period that gives the 2 max 1, obtained by a Gaussian fit to the periodogram, is assigned to the error of the true period for each TABLE 2 X-Ray Sources near GX ID Name a (J2000.0) R.A. Decl. (J2000.0) Counts s 1 (10 3 ) Separation (arcmin) 1... AX b AX c AX c AX c 19.8 Note. Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. a After Pravdo et al b SIS count rate in kev. c GIS count rate in kev.

5 No. 2, 2002 BROADENING OF IRON LINE OF GX observation. All these periods are fully consistent with the values obtained by Pravdo & Ghosh (2001). Figure 2 shows the background-subtracted X-ray pulse profile of GX obtained by folding the light curves at the best period of each observation. The pulse phase is adjusted so that the minimum intensity corresponds to phase 0.5. The midtime of this phase in each observation is listed in Table 1. As seen from Figure 2, the pulse profiles of GX have a double-peak structure, which is common among a number of X-ray pulsars (White et al. 1983), and the relative intensities of the main peak to the subpeak change with X-ray energy. Such changes of the pulse profiles have already been reported (Orlandini et al for GX 301 2). The ordinates of all the panels are plotted from zero to twice the average intensity. We can thus directly compare the pulse amplitudes shown in the different panels. In gen- (a) Apastron (b) Intermediate keV keV keV keV keV keV keV keV (c) Periastron keV keV keV keV Fig. 2. X-ray pulse profile of GX 301 2at(a) apastron, (b) intermediate, and (c) periastron phases obtained by folding the light curves at the best periods. Two phase cycles are shown for clarity. The pulse profiles in the bands of , , , and kev are arranged from top to bottom. The ordinate of all the panels begins with zero, and the average count rate is assigned in the middle.

6 884 ENDO ET AL. Vol. 574 Fig. 3. Pulse fraction of GX at the three orbital phases. For the definition of the pulse fraction, see x 3.1. Note that not all the data points are independent, but they have some overlap in the energy. eral, the pulse fraction is smaller at lower energies, and that of the periastron observation seems to be smaller than those of the other two observations. To show these characteristics quantitatively, we calculate the pulse fraction of GX for all three observations in Figure 3. The folded light curve P(E) is subject to Fourier transformation as PðEÞ ¼ C 0ðEÞ þ X1 C n ðeþ sin 2n þ n ðeþ : ð2þ 2 P n¼1 spin Using coefficients C n, the fraction of X-ray pulsation f(e)is defined by f ðeþ ¼ qffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi C 1 ðeþ 2 þ C 2 ðeþ 2 : C 0 ðeþ ð3þ We here include Fourier amplitudes C n up to the n ¼ 2 term because the folded light curve is doubly peaked. The pulse fraction falls rapidly around 6.4 kev corresponding to the iron K emission line, as indicated by Tashiro et al. (1991) Pulse Phase averaged Spectral Analysis Spectral Variation with the Orbital Phase According to the data selection criteria discussed in x 2, we extracted pulse phase averaged spectra of each detector in each observation separately. We added the SIS-0 and SIS-1 spectra into a single SIS energy spectrum for each observation. In the same way, GIS-2 and GIS-3 spectra were combined into a single GIS energy spectrum. In Figure 4 we show the background-subtracted energy spectra of the SIS without deconvolving the detector response at the three orbital phases. The spectral differences among the different orbital phases are clearly seen. In general, they fall rapidly at low energies, which is attributed to photoelectric absorption. A strong K emission line at 6.4 kev and a K edge at 7.1 kev of nearly neutral iron are clearly seen. Both of these features are produced in cold material surrounding the neutron star. The overlaid solid histograms on the three Fig. 4. Pulse-averaged energy spectra of GX at the three orbital phases. Histograms represent the best-fit model above 5 kev, composed of the power law plus Gaussian undergoing photoelectric absorption expressed by a single N H. spectra represent the model consisting of a power law plus a Gaussian with photoelectric absorption that fits the data in the band 5 10 kev. The histograms deviate from the data points below 5 kev, indicating that the single hydrogen column density is too simple to represent the spectra in the full band (Saraswat et al. 1996). The strength of the emission line and photoelectric absorption varies with orbital phase. In order to show this visually, we calculate the channel-to-channel spectral ratios of the three SIS spectra in Figure 5 with the denominator being the least absorbed apastron spectrum. In general, absorption is stronger when the separation of the two component stars is smaller. The intensity of the iron K line correlates with the absorption strength. In addition to the K line, an iron K line probably exists at least in the periastron spectrum Evaluation of the Pulse Phase averaged Spectra In order to evaluate the spectra quantitatively, we carry out spectral fits with the XSPEC, Version 11.0 software Fig. 5. PHA ratio of GX Open squares show a ratio of intermediate to apastron, and filled triangles show a ratio of periastron to apastron.

7 No. 2, 2002 BROADENING OF IRON LINE OF GX (Arnaud 1996). For the continuum, we adopt a dual powerlaw model, which consists of two power-law continua with a common photon index but with different absorbing hydrogen column densities (so-called partial-covering absorber model), as applied by Saraswat et al. (1996) to the ASCA intermediate phase spectrum. We set the abundance of iron in the absorbing matter to zero and, in turn, multiplied the iron edge at around 7.1 kev over the entire continuum model, in order to investigate the ionization state of iron in the absorbing material (Leahy et al. 1989b). For the emission line, we introduce a double-gaussian model representing iron K and K emission lines, as indicated in Figure 5. We carry out a combined fit to the GIS and SIS spectra for each observation using the spectral model described above. The results are summarized in Table 3. As expected from the discussion in x 3.2.1, the model provides acceptable fits to the apastron and the intermediate spectra, which are displayed in Figures 6a and 6b, respectively. The most notable feature of the spectra of GX is the large absorption hydrogen column density. Both column densities in the model exceed cm 2 for the spectra of all of the observations. In the case of the apastron and intermediate spectra, they are N H cm 2 and N H cm 2, whereas the covering fraction of the more absorbed power law [=Norm 1 =ðnorm 1 þ Norm 2 Þ] is larger for the intermediate spectrum, which produces a steeper cutoff at lower energies. In the periastron spectrum, both column densities are larger than those in the other two, and N H1 exceeds cm 2. We have checked the significance of the iron K line by removing it from the spectral model. The resultant 2 / values are 439/160, 286/160, and 902/156 for the apastron, intermediate, and periastron spectra, respectively, which are far larger than the values in Table 3, indicating high significance of the K line. From the errors of the Gaussian normalizations, the K line is significant at 4 level in the intermediate phase and more than 5 in the other two phases. It is important to note that K and K emission lines are both broad with a of and ev, respectively. This is much larger than what is expected from the terminal velocity of the stellar wind (x 1). We will discuss this point in detail in x 4. The 2 value of the periastron spectra indicates that the fit is formally not acceptable. By inspecting the fit residuals of the periastron spectra in detail, we find that the single- Gaussian model is too simple to fit the emission lines, in particular, the K line. We thus attempt to add two Gaussians, one for each K and K, with the width being fixed to zero. The line central energies of the newly introduced narrow components are constrained to vary with the original broad K and K components. Hence, the number of model parameters increases by two: the normalizations of the narrow K and K lines. We apply this composite line model to the periastron spectrum. The best-fit parameters are listed in the last row of Table 3, and the results are shown in Figure 6c. The resulting 2 = ¼ 171:0=151 implies that the fit is marginally acceptable. We also perform F-tests to check the significance of the two additional line parameters. As a result, the introduction of the composite line model is justified at the level of more than 99.9%. We conclude that the iron K emission line profile in the periastron spectrum can be represented by the narrow core component and the broad wing component that are both represented by Gaussians with ¼ 0 and 135 ev, respectively. For the K line, on the other hand, the introduction of the narrow core component is not significant, and only an upper limit of the intensity is obtained. We note, however, that the composite line model simply implies that the iron K line profile is too complicated to be represented by the single Gaussian. Its physical implication will be discussed in x 4. Based on the results of the spectral fit, we summarize the observed flux and the bolometric luminosity in Table 4. In calculating the bolometric luminosity, we account for the exponential cutoff at higher energies. Since the cutoff starts above 20 kev (x 1), which is outside the ASCA bandpass, we are not able to determine the cutoff parameters. We thus refer to Orlandini et al. (2000), who derived both E c and E f at various orbital phases using the BeppoSAX data. The distance to GX is assumed to be 1.8 kpc (Parkes et al. 1980). The bolometric X-ray luminosities of the intermediate and the apastron phases are nearly identical at ergs s 1, whereas that at the periastron phase, ergs s 1, is larger than at the other two phases by roughly an order of magnitude Confirmation of the Broad Line Profile It is known that the performance of the SIS has been degraded since launch (1993 February 20) as a result of accumulation of radiation damage in space (Yamashita et al. 1999), and its effect on the energy-resolving power cannot be neglected at the times of the apastron and the periastron observations of GX Hence, we check whether its effect is correctly reflected on the energy response function of the SIS. For this purpose, we attempt to apply the composite iron line profile of GX to that of an X-ray source, which has a narrow iron emission line in its spectrum and was observed almost at the same time as the apastron and periastron observations. By surveying the ASCA archival data, we find that the Centaurus Cluster was observed from 1995 July 19.4 to 21.7, only half a year before the periastron observation of GX The spectrum of the Centaurus Cluster can be represented by optically thin thermal plasma emission with a temperature of 4 kev, including a narrow He-like iron K emission line centered at 6.70 kev (Ikebe et al. 1999). The net exposure time was 72.0 ks for GIS and 69.3 ks for SIS. During this observation, the GIS was operated in the PH mode with the nominal telemetry bits configuration. The SIS was operated in the 1-CCD faint mode throughout. We have screened the Centaurus Cluster data as described in x 2.2. In the analysis of Centaurus Cluster data, we concentrate on the SIS data. The faint mode energy spectrum of the SIS is accumulated from the entire area of the CCD chip. The background spectrum was made from the archival database of blank sky observations. We then fit the background-subtracted SIS spectrum above 5 kev with the optically thin thermal plasma emission model (MEKAL) in XSPEC (Mewe, Gronenschild, & van den Oord 1985; Mewe, Lemen, & van den Oord 1986; Kaastra 1992; Liedahl, Osterheld, & Goldstein 1995) with the redshift being fixed at the optical value of z ¼ 0:0104 (Ikebe et al. 1999). We obtain a plasma temperature of 4:29 0:45 kev and elemental abundances of 0:72 0:06 Z with a 2 value of 2 =dof ¼ 22:9=30, where we have adopted the element composition table by Anders & Grevesse (1989) in the fitting

8 TABLE 3 Pulse Phase averaged Spectral Parameters of GX A. Continuum Phase 2 /dof C Norm 1 (photons cm 2 s 1 kev 1 ) Norm 2 (photons cm 2 s 1 kev 1 ) N H1 (10 23 cm 2 ) N H2 (10 23 cm 2 ) E edge (kev) edge Single Gaussian for Each K and K Apastron /157 1:29 0:03 þ0: :41 þ3:6 1:44 2:02 þ0:03 0:07 7:18 þ0:11 0:27 0:27 þ0:10 0:07 Intermediate /157 1:59 þ0:12 Periastron /153 1:34 þ0:13 0:11 0:34 þ0:12 0:08 0:02 þ0:02 0:01 7:09 þ0:68 0:14 1:92 þ0:85 0:62 0:20 þ0:12 0:10 11:5 1:0 4:69 þ0:47 0:58 2:25 þ0:52 0:58 7:15 þ0:03 0:06 0:80 þ0:11 0:10 0:53 7:16 þ0:03 0:04 1:21 þ0:13 0:11 Double Gaussian for Each K and K Periastron /151 1:18 0:02 þ0:12 1:25 þ0:55 0:30 0:12 þ0:06 0:02 10:5 0:2 þ0:3 4:34 þ0:23 0: B. Emission Lines Phase E K (kev) K B/N (ev) I K B/N (10 3 photons cm 2 s 1 ) EW K B/N (ev) E K (kev) K B/N (ev) I K B/N (10 3 photons cm 2 s 1 ) EW K B/N (ev) Single Gaussian for Each K and K Apastron þ26 29 Intermediate þ16 23 / / / <100/ /... 42þ13 12 /... /... 4:6þ0:5 0:4 Periastron... 6:400 0:001 þ0: / þ28 20 / þ26 20 / þ / þ110 7:06þ0:08 0:07 69 / þ23 25 /... 0:7þ0:4 0:3 /... 35þ9 7 /... 44þ24 19 /... / þ84 64 /... Double Gaussian for Each K and K Periastron... 6:395 0:001 þ0: þ11 15 /0 (fix) 499 þ63 7:14 þ0:02 80 þ45 76þ10 4 /68þ /443þ :01 26 /0 (fix) 34 5/< þ40 38 /<24 Note. B/N denotes broad/narrow components.

9 BROADENING OF IRON LINE OF GX (a) Apastron GIS SIS (b) Intermediate GIS SIS (c) Periastron GIS SIS Fig. 6. Average SIS spectra of the (a) apastron, (b) intermediate, and (c) periastron phases fitted by the dual absorbed power laws with the common photon index and iron K and K lines. The K line can be fitted by the single broad Gaussian with ¼ ev in the apastron and intermediate spectra. On the other hand, narrow ( ¼ 0) and broad ( ¼ 135 ev) double Gaussians are necessary to fit the K line in the periastron spectra. with XSPEC. These results are consistent with those of Ikebe et al. (1999). The model indicates the existence of Helike iron K and K lines and those from H-like iron. We then replace the MEKAL model with a thermal bremsstrahlung plus four Gaussians, representing the iron lines, and see whether the width of the emission lines can be derived correctly. Before doing this, however, we have transformed the data format from original faint mode to bright mode, since the spectral analysis of GX has been carried out in bright mode. The best-fit SIS-0 þ SIS-1 spectrum in the range of 5 10 kev is shown in Figure 7a. The 2 value for this fit is 2 =dof ¼ 24:3=27. The resulting bremsstrahlung temperature is found to be 4:27 0:66 kev. The best-fit width of the He-like iron K line is consistent

10 888 ENDO ET AL. Vol. 574 Phase TABLE 4 Observed Flux and Bolometric Luminosity Observed Flux a (10 10 ergs cm 2 s 1 ) Bolometric Luminosity b (10 36 ergs s 1 ) Apastron... 6:85 þ0:05 0:10 2:4 0:3 þ0:2c Intermediate.. 2:91 þ0:08 0:07 2:1 þ0:6 0:5 d Periastron e a In the 2 10 kev band, absorption is not removed. b Distance is assumed to be 1.8 kpc; absorption is removed. c E c ¼ 22:0 kev and E f ¼ 11 kev (Orlandini et al. 2000). d E c ¼ 22:5 kev and E f ¼ 12 kev (Orlandini et al. 2000). e E c ¼ 19:5 kev and E f ¼ 8 kev (Orlandini et al. 2000). with zero, with an upper limit of ¼ 57 ev. This result demonstrates that the effect of the radiation damage is properly reflected to the energy response function. We then fit the spectrum of the Centaurus Cluster with the same model but with the He-like iron K line profile replaced by that obtained from the periastron spectrum of GX (Fig. 6c; Table 3). We fixed the values of the double Gaussians at 0 and 135 ev and their intensity ratio at 0.89 (Table 3). The best-fit result is shown in Figure 7b. We obtain 2 / of 40.3/27, which indicates that the model is not acceptable at the 90% confidence level. Therefore, the SIS can discriminate the composite line profile found in GX spectrum from the narrow line, even at the time of the periastron observation. We conclude that the composite line profile of the iron K emission line in the periastron spectrum is real and is not due to the degradation of the energy resolution of the SIS Pulse Phase resolved Spectral Analysis In order to investigate variations of the composite iron K line profile, we perform a pulse phase resolved spectral analysis of the periastron data. Because of statistical limitations, we can sort the photons into only four evenly spaced phase bins, such as spin ¼ 0:00 0:25, , , and (note that spin ¼ 0:5 corresponds to the intensity minimum of the folded light curve; Fig. 2). We thus continue to shift the phase boundaries stepped by 0.1. As a result, we obtain five sets of data composed of four independent spectra (total of 20 spectra). These 20 phase-resolved spectra are individually fitted above 2 kev with the same model as the phase-averaged spectra (x 3.2.2). The phase-averaged background spectrum is used for the background subtraction before the fitting. We fix the central energy of the narrow iron K line and the energy of the iron K edge to the values obtained by the fit to the phase-averaged spectra (last row of Table 3) because the scatter of the best-fit values of these parameters obtained from 20 spectra is less than 1%, which is well within the statistical errors. We plot some of the best-fit parameters in Figure 8: (a) the power-law photon index, (b) the absorption-corrected continuum luminosity in 2 10 kev, (c) the width of the broad wing component of iron K line with Gaussian,(d ) the central energy of the broad wing component of the iron K line, and (e) the sum of the intensities of the broad and narrow iron K components. The central energy of the broad wing is nearly constant at 6.40 kev throughout the pulse cycle. The line intensities of the broad wing and narrow core components are strongly coupled with each other in the spectral fitting. We therefore plot the sum of the two. The photon index appears to be constant throughout the spin cycle, which is in contrast to the results so far reported (see, e.g., Tashiro et al. 1991, in which C varies in the range with Ginga). A possible interpretation for this difference is the ASCA bandpass being limited only up to 10 kev, much lower than that of Ginga. This makes the ASCA observations free from the Compton component (the positive power-law component described in x 1; see Makishima et al. 1999). The difference between the two results can thus (a) Simple Gaussian Model χ 2 /d.o.f.=24.3/27 Centaurus cluster (b) Composite Line Model χ 2 /d.o.f.=40.3/27 Fig. 7. Bright mode SIS energy spectrum of Centaurus Cluster in the 5 10 kev band fitted by (a) thermal bremsstrahlung plus narrow He-like iron K line and (b) thermal bremsstrahlung plus the composite line profile for the He-like iron K line. The latter model does not fit the spectrum.

11 No. 2, 2002 BROADENING OF IRON LINE OF GX Fig. 8. Pulse phase dependence of the spectral parameters in the periastron phase: (a) power-law photon index, (b) 2 10 kev absorption-corrected continuum luminosity (without line component) in units of ergs s 1,(c) width of the broad wing component in Gaussian (kev), (d ) central energy of the broad wing component (kev), and (e) sum of the broad wing and narrow core intensities in units of photons cm 2 s 1. All the errors are 90% confidence level. be understood if the photon index of the intrinsic power law is stable while only the unsaturated Compton component depends on the spin phase. Further investigation is, however, necessary to understand the apparent discrepancy between the Ginga and ASCA results. The width of the broad wing component in panel (c) seems to vary significantly with the spin phase. It shows a doubly peaked modulation with the peaks and valleys corresponding to the valleys and peaks of the continuum luminosity, respectively. Among the five sets of data, those displayed by open circles exhibit the largest variation in line width, and the significance of the variation between phases N( spin ¼ 1:20 1:45) and B ( spin ¼ 0:95 1:20) and between phases N and B 0 ( spin ¼ 1:45 1:70) is 3.7 and 5.1, respectively. The width of the broad wing, however, is possibly coupled with its intensity. Thus, we attempt to draw confidence contours in the plane of the intensity and the width of the broad wing component. Before doing this, we have summed up the spectra of phases B and B 0 to increase statistics, to be compared with phase N. The result is shown in Figure 9. The dashed, solid, and dotted contours represent 68%, 90%, and 99% confidence levels, respectively. Among them, the 90% confidence contours do not overlap with each other. From this analysis, we conclude that the profile of the broad wing component varies between phase B þ B 0 and phase N. 4. DISCUSSION 4.1. Emission Site of Iron K Line Implication of the Line Broadening It has been believed that the iron K emission line from GX is emitted from the stellar wind spreading widely over the binary orbit scale. Current ASCA results on the line width, however, pose a question to this picture from an observational point of view. The terminal velocity v 1 of the stellar wind from WRA 977 is estimated to be 400 km s 1 (Parkes et al. 1980). This amounts to only v 1 =c ¼ 1: and can cause a line broadening of only 8.5 ev at

12 890 ENDO ET AL. Vol. 574 exceed cm, which also indicates that the iron emission line is formed in a region within R acc. After passing through R acc, matter starts to accrete along the magnetic field lines at the Alfvén radius r A defined by Fig. 9. Confidence contours of the intensity vs. the width of the broad wing component in the periastron spectra. The crosses are the best-fit values. Dashed, solid, and dotted lines represent contours of 1, 90%, and 99% confidence levels, respectively. the iron K line energy. On the other hand, as shown in x 3.2, the iron K line during the apastron and intermediate phases is represented by the broad Gaussian with the width of ¼ 40 þ16 23 and 72þ26 29 ev, respectively. For the periastron phase, the width of ¼ 79 þ6 6 ev is obtained if we model the line by the single Gaussian (although the fit is not acceptable as shown in Table 3), and ¼ 135 þ11 15 ev if we adopt the composite line model (Table 3). All of these values are larger than that expected from the terminal velocity by roughly an order of magnitude. To explain the observed line widths, fluorescence must occur at a much higher velocity region, which can be realized only by the gravity of the neutron star. The iron line emission site is undoubtedly limited to a space much closer to the neutron star than ever imagined. Material in the stellar wind begins to accrete toward the neutron star at the accretion radius R acc ¼ 2GMv 2 rel (x 1). In the case of GX 301 2, R acc ¼ 2GM v 2 2GM rel v 2 ¼ 2: M 1 1:4 M v km s 1 ðcmþ : ð4þ It is noteworthy here that Tashiro et al. (1991) derived, based on the Ginga observation, that the distance of the iron line emission site from the neutron star does not greatly 1 2 ðr AÞv 2 ðr A Þ¼ B2 ðr A Þ ; ð5þ 8 where and v are the mass density and the velocity of the accreting matter, respectively. Assuming that v can be approximated by the free-fall velocity and the dipole magnetic field (B ¼ l=r 3 ), and transforming by _M in spherical accretion with _M ¼ L X R=GM, we obtain r A ¼ 3: l 4=7 30 M1=7 1:4 L 2=7 37 R 2=7 6 ðcmþ ; ð6þ where M 1.4 is M/1.4 M, l 30 ¼ l=10 30 Gcm 3, and so on. Note that l 30 ¼ 3:2 (Makishima & Mihara 1992; Mihara 1995; Makishima et al. 1999) if R 6 ¼ 1. The free-fall velocity at r A is v ff ðr A Þ c ¼ 1 c sffiffiffiffiffiffiffiffiffiffiffiffiffiffi 2GM X r A ¼ 0:033l 2=7 30 M 3=7 1:4 L1=7 37 R1=7 6 : ð7þ In Table 5 we summarize the estimates of r A and v ff ðr A Þ based on equations (6) and (7), assuming l 30 ¼ 3:2 and M 1:4 ¼ R 6 ¼ 1, for comparisons with the observed energy widths obtained from the three ASCA observations. The observed energy widths evaluated based on the single-gaussian model are in the range 1 3 of those expected from the free-fall velocity at r A for the three observations. This result indicates that the iron K line is emitted from a region that is (4 9)r A from the neutron star, or closer since the free-fall velocity is the maximum possible velocity at a given radius. We conclude that the iron K line of GX originates from a region within 10r A (10 10 cm) of the neutron star. The size of the iron K line emission site in HMXBs has been estimated by utilizing the eclipse of the neutron star by the optical companion star. For Vela X-1, Ohashi et al. (1984) found from Tenma observations that the intensity of the fluorescent iron K line during mideclipse phase is reduced to 6% of that out of the eclipse and concluded that the size of the line emission site is within cm from the neutron star, which is smaller than the binary separation a ¼ 53:4 R or cm (van Kerkwijk et al. 1995) by roughly an order of magnitude. In a similar way, Ebisawa et al. (1996) estimated the size to be less than cm for Cen X-3. Nagase et al. (1992) gave a much tighter constraint of cm for Cen X-3 by analyzing a series of Ginga spectra (although Cen X-3 is believed to be a disk-fed pulsar and reflection from the accretion disk is another candidate 1 2 TABLE 5 Fractional Energy Width of Iron K Line and That Expected at r A Orbital Phase Luminosity (10 36 ergs s 1 ) (ev) /E K (10 2 ) r A (10 8 cm) v ff ðr A Þ=c (10 2 ) Apastron þ :13 þ0:41 Intermediate þ :63 þ0:11 Periastron þ6 6 1:23 þ0:09 Periastron a þ :10 0:23 þ0:17 a Broad wing component in the composite line model. 0: : :

13 No. 2, 2002 BROADENING OF IRON LINE OF GX of the iron K emission). Although only a few examples have been found as yet, we infer that the iron K emission line from HMXB pulsars generally originates from a region well within the accretion radius R acc, close to the Alfvén radius Reviewing the Broad Wing Component in Periastron In x 3.2 we have pointed out that the composite line profile is necessary to fit the iron K emission line in the periastron spectrum, and in x 3.3 we have shown that its shape is modulated according to the spin phase. Since the mass accretion geometry out to the Alfvén radius is spherically symmetric, it is necessary to invoke accretion columns to account for the spin phase dependence. We propose a scenario in which the broad wing component of the iron K line originates from preshock accretion columns irradiated by the pulsar. In this scenario, K photons generated in the approaching and in the receding columns are blue- and redshifted, respectively. The degree of the energy shift depends on the angle between the dipole axis and the observer s line of sight, which causes the spin modulation of the line profile. In order to realize this idea in the spectral fit, we substitute double Gaussians corresponding to the fluorescent blue- and redshifted lines from the dual columns for the broad Gaussian hitherto adopted for the broad wing component. The central energies of the newly introduced double Gaussians are determined by a single parameter representing the velocity of matter in the accretion columns (=v=c)as ð1 ÞE K, where E K is the energy of the iron K line in the rest frame (6.40 kev). For the K line, on the other hand, we have simply adopted a single broad Gaussian because we have only an upper limit for the narrow Gaussian in x 3.2 (Table 3). First, we assume that the width of the core component of the K line is narrow ( is fixed at 0 ev), as in the composite line model, and fit the spectra of phase N ( spin ¼ 0:20 0:45) and phase B ( spin ¼ 0:95 0:20) simultaneously using the dual power-law model (x 3.2) plus iron emission lines described above in the 5 10 kev region. The result of the fit is shown in Figure 10a (the band kev is expanded), and the best-fit iron K line parameters are summarized in Table 6. The reduced 2 value is 0.77, and the fit is acceptable. The resulting is and for phases N and B, respectively, which are almost equal to or larger than that expected from the free-fall velocity at the Alfvén radius (Table 5). We next set the width of the core component free to vary but constrain them to be common between the two phases. The result of the fit is displayed in Figure 10b, and the bestfit parameters are listed in Table 6. The best-fit width of the core component has a nonzero value of 64 þ16 31 ev, which is comparable with that obtained from the apastron and periastron spectra. The intensities of the blue- and redshifted components are similar between phases N and B, and the total equivalent width is ev, which is smaller than that of the core component ( ev) by roughly an order of magnitude. The resulting 2 value decreases by 3.8 from the previous fit with one additional model parameter. F-test indicates that the significance of adding the parameter is 97.4%. Although it is difficult to see the improvement of the fit from Figure 10a to Figure 10b, we believe that the latter broad core picture is more realistic because the iron K lines in the apastron and intermediate phases are both broad. These emission lines originate from the region within 10r A (x 4.1.1), and it is natural to ascribe the core component in the periastron phase also to this region. Note that we have assumed that the blue- and redshifted components appear symmetrically with respect to the energy E K ¼ 6:40 kev. There are, however, two possible mechanisms that may invalidate this assumption: gravitational effects of the neutron star and the presence of magnetic fields. First, as a result of strong gravity, the energy of the emission lines from the accretion column can be significantly shifted from that in the laboratory. In a Schwarzschild metric, the fractional energy shift is estimated to be E obs E K ¼ s ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi 1 2GM rc 2 sffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi ¼ 1 0:41 M 1:4 ; ð8þ r 6 which indicates that the iron K line should appear at E obs ¼ 4:9 kev if it is emitted from the neutron star surface. As will be shown in xx 4.2 and 4.3, however, the iron K emission in GX is attributed to a region with r > 10 7 cm. In that case, equation (8) indicates that the energy shift is only 130 ev. As can be understood from Figure 10, it is difficult to resolve an energy shift at this level, as a result of the existence of the broad core component and the limited energy resolution of the SIS. Second, the Zeeman effect due to strong magnetic fields is also possible to alter the E K significantly from 6.40 kev. However, the magnetic field is not much stronger than 10 9 Gatr > 10 7 cm. The corresponding splitting energy under this condition is estimated to be DE ¼ eh 4m e c B ¼ 5:8 B 10 9 G ðevþ : ð9þ This is much smaller than equation (8) and can be neglected in the present case. In summary, we conclude that the core component of the iron K line exists at 6.40 kev at all orbital phases with a similar width of ¼ ev. In addition, the broad wing component found only in the periastron phase can be interpreted as the blue- and redshifted fluorescent iron K components from the bipolar accretion columns Other Possibilities for the Broad Wing Component In the previous subsection, we introduced a scenario in which fluorescence lines are produced in preshock accretion columns to explain the broad wing component of the iron K line without any detailed justifications. In this subsection, we consider another possibility, Compton scattering in the accretion columns, to account for the broad wing component. If the seed 6.4 kev photons with spatially uniform distribution are injected into the accretion columns, the fractional shift of the energies scattered off the receding and the approaching columns of order 1 and 1 þ, respectively, is expected. Although this picture can naturally explain the apparent symmetry of the broad wing component centered on 6.4 kev (Fig. 8), it is difficult to find a source of the 6.4 kev seed photons strong enough to explain the observed equivalent widths ( ev; Table 6). The primary candidate of the seed photon is the core component with ¼ 64 ev and an equivalent width of ev. By comparing the equivalent widths, however, the accretion