Table 5.3—Continued
Orbital phase line Energy Sigma Intensityb
(keV) (eV) (photon cm−2 s−1)
PP Cr Kα 5.41140 0.00 7.53e−5
(5.40826–5.41166) (0.00–5.98) (4.99–10.2)
PP Fe Kαa 6.39469 13.2 8.12e−3
(6.39421–6.39516) (12.5–13.9) (7.94–8.31)
PP Fe Kβ 7.05755 13.8 1.54e−3
(7.05476–7.06035) (9.85–17.5) (1.37–1.69)
PP Ni Kα 7.45984 0.00 5.07e−4
(7.45682–7.46916) (0.00–15.5) (3.73–6.40)
aThe Compton shoulder region (6.24–6.34 keV) is excluded.
bInter stellar gas absorption is corrected. The hydrogen column density of 6 × 1021 cm−2 is assumed, corresponding to the density of 1 H cm−3 and a distance of 1.8 kpc.
Note. — Errors correspond to 90 % confidence levels.
Figure 5.5: The blown-up spectra of the iron Kα lines in the PP (bottom), the IM (middle) and the NA (top). A shoulder component extending toward the low-energy side of the line is clearly seen and fully resolved. The width of the shoulder (∆E ∼160 eV: 6.24–6.40 keV) matches the energy distribution of iron Kα photons that suffer single Compton scattering. The line shows the best-fit model in the spectral fits of the continuum.
Figure 5.6: Pulse profiles of GX 301−2 in IM (top panel), NA (middle) and PP (bottom) by folding the light curves at the periods of 683.2 s, 680.8 s and 681.2 s, respectively. Two phase cycles are shown for clarity. The HEG± 1 order events are used, and the energy band is 1–10 keV. The pulse phase of 0.0 corresponds to the start time of each observation.
Table 5.4. Total and iron Kα fluxes of pulse phase divided spectra
orbital Flux (2.0–10.0 keV)(erg s−1) Fe Kα flux (photon cm−2 s−1)
phase peak(1) peak(2) peak(1) peak(2)
IM (2.5 ± 0.1) ×10−10 (1.8 ± 0.1) × 10−10 (8.8 ± 1.1) × 10−4 (6.9 ± 0.8) × 10−4 NA (1.5 ± 0.2) × 10−9 (1.1 ± 0.1) ×10−9 (3.3 ± 0.2) × 10−3 (3.2 ± 0.2) × 10−3 PP (1.2 ± 0.1) × 10−9 (8.9 ± 0.2) × 10−10 (7.8 ± 0.4) × 10−3 (7.7 ± 0.4) × 10−3
Figure 5.7: X-ray spectra divided by pulse phase. The magenta spectra are obtained from the events in peak(1), whose count rates are relatively high (see Figure 5.6). The cyan spectra are extracted from the events in peak(2).
Chapter 6
Simulation of Photoionized Plasma in HMXB
6.1 Modeling of Photoionized Plasmas
The unprecedented spectral resolution of the grating system onboard the Chandra satel-lite gives us a wealth of information about X-ray emission lines from HMXBs, Vela X-1 and GX 301−2. As already shown in previous chapters, we have succeeded in detecting lines from highly ionized ions, such as H-like and He-like Si, S, Mg, from photoionized plasma, together with clear signals from radiative recombination continua in Vela X-1.
As expected from the dynamical motion of the gas in the stellar wind, opposite shifts of the central energy appear to exist in phase 0.50 and eclipse for Vela X-1. On the other hand, GX 301−2 shows only fluorescent lines from low charge states. From the detail analysis of the line profiles of the iron fluorescence line, we discovered a shoulder-like structure which could be accounted by the effect of Compton scattering.
The emission lines we detected from Vela X-1 and GX 301−2 can be interpreted as due to emission from a gas photoionized by the X-ray radiation from the neutron star. In this case, the spectrum emerged from the binary system is the result of the propagation of X-rays through the gas. Line emission, presumably due to processes, such as photoionization, recombination and fluorescence, are controlled by the ionization structure and the density distribution of the gas in the stellar wind. Therefore, by investigating characteristics of lines, we will be able to obtain an important clue to addressing the questions about how the photoionized plasma is distributed spatially in the stellar wind and how the distributions can affect the nature of the X-ray emission observed in HMXBs.
In previous studies of photoionization plasmas, modeling were performed based on very simple assumptions. Kallman & McCray (1982) presented theoretical models to calculate the ion abundances and the temperature for the given X-ray spectrum and the wind density. Since the introduction of their model as a computer code (called XSTAR
1), it has been frequently used to characterize X-ray spectra from HMXBs. However, most of the analysis are based on very simple assumptions such as symmetric geometry and constant density distribution in the volume.
In order to find the ionization structure from the HETGS results, it is necessary to model the emission from more realistic environments in the stellar wind. Therefore, we have developed a new code to simulate X-ray interactions in the stellar wind. Most important ingredients in the code is that fully three dimensional treatment is adopted for both the ionization structure and the photon transportation in the plasma. For the case that the optically thin approximations are allowed, the condition of photoionized plasma can be characterized only by the ionization parameter ξ. However, the Chandra results indeed imply that the emission lines with high ionization degree actually come from a dense region (particle density n >109 cm−3). In this situation, we have to consider not only the ξs, but also the effect of absorption by the matter between the neutron star and the emission site, because, the incident flux is absorbed by the material and changes the shape and the flux.
Another important issue to be stressed is that we now see the effect of the Doppler shift due to the emission from the material moving with fast velocity. The information on the Doppler shift could also constrain the velocity field of the X-ray emitting region in the stellar wind. The model to be used for the re-construction of physical environment in the stellar wind has to deal with these kinematical effects.
Our simulation code consists of two parts:
(1) Calculation of the ionization structure,
(2) Monte Carlo simulation for tracking X-ray photon transportation.
In part (1), we construct the map of the ion abundance in the stellar wind. In part (2), we have implemented the physical process which are related to highly ionized gas, together with the Lorentz transformation for the calculation of Doppler effects. The simulation is held in three dimensional space divided into grids such that we can handle more complex geometries. Our procedure of each of these parts is described in the following sections.