Suzaku Investigation of Hard X-Ray Emission Associated with the Galactic Center Region

Suzaku Investigation of Hard X-ray Emission Associated with the Galactic Center Region

Ken-ichi Tamura

Department of Physics Graduate School of Science University of Tokyo

December 19, 2007 Abstract

In the center region of our Galaxy, the intense diffuse X-ray emission exists. The emission has been observed as X-ray spectra from a hot plasma with a temperature of 108 K, and the mechanism of production has been still a mystery over the past two decade from the discovery. The X-ray observatory, Suzaku, carries the X-ray CCD cameras, XIS, with good energy resolutions and the Hard X-ray Detector, HXD, with high energy band up to several hundreds keV, and therefore, has a large advantage to solve the mystery. Especially, with HXD-PIN which is a component of HXD , we will be able to obtain hard X-ray spectra of the Galactic center diffuse emission if we can estimate the contamination from the hard X-ray sources in the FOV. Hence, we have performed the in-orbit calibrations of the angular response and developed the new method to estimate the contamination. We have performed the mapping observations for the Galactic center region with 35 pointing and the total exposure of 1 Msec. Both XIS and HXD-PIN have detected strong X-ray emissions in all the observations. From the detailed XIS analysis, we have shown that the shapes of the XIS spectra are uniform everywhere. And then, we have confirmed the distribution of the surface brightness extracted from the XIS spectra is consistent with that of Fe line in the hot plasma. For the HXD-PIN data, we have estimated the contamination from the hard X-ray source in the detected hard X-ray fluxes and shown that significant hard X-ray fluxes remain in entire the region. And then, we have analyzed the XIS and HXD-PIN wide-band spectra to show that the spectra can not be explained by only the thermal emissions and the power-law component exists certainly. Finally, we have shown that the distribution of the thermal and non-thermal emissions have strong correlation and suggested that both emissions are radiated form the same origin. Contents

1 Introduction 5

2 Review 7 2.1 Overview of the Galactic Center Region ...... 7 2.2 Past X-ray Observations of the Galactic Center Region ...... 9 2.2.1 Ginga results ...... 9 2.2.2 ASCA results ...... 10 2.2.3 Superposition of Dim Point Sources ...... 10 2.3 Non-thermal Hard Tail Associated with the Milky Way ...... 12 2.4 Suzaku results ...... 13

3 The X-Ray Observatory Suzaku 15 3.1 The Suzaku Spacecraft ...... 15 3.2 X-Ray Telescope (XRT) ...... 17 3.3 X-Ray Imaging Spectrometers (XIS) ...... 22 3.4 Hard X-Ray Detector (HXD) ...... 26 3.4.1 Overview ...... 26 3.4.2 HXD-PIN Detectors ...... 31 3.4.3 In-Orbit Calibration ...... 32

4 Angular Response of HXD-PIN and A New Method for Flux Estima- tion 34 4.1 Angular Response of HXD-PIN ...... 34 4.1.1 Fine-Collimator ...... 34 4.1.2 Angular Response ...... 35 4.2 In-Orbit Calibration of the Angular Response ...... 35 4.2.1 Calibration of Light Axes ...... 35 4.2.2 Fine Tuning of the Angular Response ...... 42 4.3 New Method for Flux Estimation ...... 43

5 Suzaku Observations 50 5.1 Overview ...... 50 5.2 Strategy of the observations ...... 50

2 5.2.1 Determination of the temperature of the hot plasma ...... 51 5.2.2 Monitoring bright transient hard X-ray sources ...... 51 5.2.3 Observations of molecular clouds ...... 52 5.2.4 Mapping observations ...... 52 5.3 Status of the Observations ...... 52

6 Suzaku Data Analysis and Results 55 6.1 XIS Data Analysis ...... 55 6.1.1 Data Reduction ...... 55 6.1.2 Imaging analysis ...... 55 6.1.3 Analysis of the XIS spectra ...... 59 6.1.4 Surface Brightness Distribution of Soft X-ray ...... 67 6.2 HXD-PIN Data Analysis ...... 70 6.2.1 Data Reduction ...... 70 6.2.2 NXB Modeling ...... 70 6.2.3 Spectral Analysis ...... 71 6.2.4 Spectral Fitting with a Power-law Model ...... 72 6.3 Estimation of Contaminations from Known Hard X-ray Sources . . . . . 85 6.3.1 Known Hard X-ray Sources in the Galactic Center Region . . . . 85 6.3.2 Flux Estimation of the Hard X-ray Sources with IBIS ...... 86 6.3.3 Spectral Estimation of the Hard X-ray Sources with XIS . . . . . 95 6.3.4 Flux Estimation of ”1E 1740.7-2942” ...... 96 6.4 Analysis of the Hard X-ray diﬀuse emission ...... 106 6.4.1 Class A ...... 106 6.4.2 Class B ...... 109 6.4.3 Class C ...... 116 6.4.4 Distribution of Hard X-ray Emission ...... 116

7 Discussion 120 7.1 Brief Summary of the Observational Results ...... 120 7.2 Uncertainties of the Hard X-ray Diﬀuse Emission ...... 120 7.3 Contributions from the Dim Point Sources ...... 122 7.4 Interpretation of the Hard X-ray Diﬀuse Emission ...... 123

8 Conclusion 125

A Image of Galactic Center Region with Swift 126

B Individual Observational Data 128 B.1 The XIS spectra and the IBIS ﬂuxes of the bright point sources . . . . . 128 B.2 The XIS spectra ﬁtted with the template model of the west region . . . . 130

3 B.3 The HXD-PIN spectra subtracted the contaminations from the known hard X-ray sources...... 136 B.4 The XIS and HXD-PIN Unfolded Spectra ...... 141

4 Chapter 1

Introduction

The intense iron line emission concentrating in the center of our Galaxy has been observed in X-ray band with previous X-ray satellites. This emission often interpreted as diffuse hot plasma with a high temperature of ∼ 6 keV being confined in the galactic center region. Despite of repeated observations and numerous efforts to solve the origin of this diffuse hot plasma, over the past two decades from the discovery of the emission, both ”How the hot plasma has been created?” and ”Why the hot plasma can stay there?” have been mysteries. In addition, diffuse X-ray emissions from the Galactic plane and bulge have been also observed. From these regions, non-thermal diffuse emissions have also been discovered. However, no conclusions of the emission mechanisms have been obtained. From energetics point of view, the origin of hot plasma confined in the Galactic Center Region is of great importance in modern high energy astronomy. Recently, a symptom of the non-thermal emission has been reporeted from the analysis of the XIS data, (e.g. Koyama et al. 2007), which is the CCD sensor onboard Suzaku. Suzaku is the fifth Japanese X-ray observatory which has two detector systems. One is XIS, instruments with X-ray CCD cameras, which covers the energy range of 0.2–12 keV. The other detector, HXD, is a collimated detector which features the narrowest field- of-view and the lowest background among recent collimator-type hard X-ray detectors. These features give large advantages to observations for the Galactic center region in which many hard X-ray sources, giant molecular cloud and possible the Galactic center diffuse emission are mixed with each other. The objective of this thesis is to establish the existence of non-thermal X-ray diffuse emission in the galactic center region, and to show how the thermal and the non-thermal emissions are correlated. In Chapter 2, we review the previous works and the current understanding on the Galactic diffuse emission. Chapter 3 gives a brief description of the Suzaku satellite and its detector systems used for our observations. In Chapter 4, we explain the in-orbit calibrations of the angular response which is essential for the region containing a lot of hard X-ray sources. In Chapter 5, the overview and strategy on the observations of the Galactic center region are described. In Chapter 6, we show the

5 analysis and result of the XIS and HXD data, and the estimation of the contaminations from the hard X-ray sources to the HXD ﬂuxes. In Chapter 7, we discuss the Suzaku results. Finally, conclusions of this thesis is given in Chapter 8.

6 Chapter 2

Review

2.1 Overview of the Galactic Center Region

The Galactic center is the closest galactic nucleus of a spiral galaxy, and hence numerous observations have been performed on this region. The central half kiloparsec region around the Galactic Center is an extremely complex region containing a variety of astro- physical activities: cold and warm molecular clouds, star cluster/formation, supernova remnants (SNRs), and HII regions, to name a few. Therefore, this region have been attracting many researchers in wide-band wavelengths, from radio to very high energy gamma-ray.

Figure 2.1: Velocity integrated (-200 to 200 km s−1) CS J=1-0 emission in the Galactic center region, obtained with 45 m telescope at Nobeyama Radio Observatory (Tsuboi et al. 1999). The data have been convolved with a 60” Gaussian.

The structures of the molecular clouds, the fountainhead of star formation activities and hot plasmas eventuated, have been mapped with CO and CS lines in the radio wavelength. As shown in ﬁgure 2.1, it is clear that they show a strong concentration around the center of Galaxy as a form of giant molecular clouds involved in a continuous ridge extended along the Galactic plane. The compact and luminous nuclear region produces ∼5–10% of our Galaxy’s infrared and Lyman continuum luminosities and accounts for

7 Figure 2.2: A mosaic image obtained with a survey of Spitzer Space Tele- scope/IRAC observations of the central 2 × 1.5 degrees (265 × 200 pc) of the Galaxy at 3-8 µm. rougly 10% of the molecular gas content. Overlaid on this global structure, some peculiar structures such as the expanding molecular ring, a center lobe which are often explained as a remnant of huge explosive events, and the radio arc and plumes possibly resulted as a local consolidation of the magnetic field, have been observed (Kaifu et al. 1972, Yusef-Zadeh et al. 1984, Tsuboi et al. 1985). The highest spatial resolution and sensitivity large-scale map made to date of the Galactic Center at mid-infrared wavelengths has been recently obtained with the Spitzer telescope as shown in figure 2.2 (Stolovy et al. 2006). Since the opacity of interstellar medium is small in the infrared wavelength, these oservations provide a census of the optically obscured stellar sources, in addition to a detailed map of the highly flamentary structure in the interstellar medium. At this wavelength, so far no counterparts to the unique system of nonthermal radio filaments present has been found.

8 2.2 Past X-ray Observations of the Galactic Center Region

2.2.1 Ginga results

Based on the discovery of the Fe-K line emission by Tenma (Koyama et al. 1986), an extensive survey of the Fe-K line was performed along the Galactic plane using the Large Area Counter (LAC) aboard Ginga, with a filed-of-view of 1◦.0×2◦.0 FWHM (Yamauchi et al. 1990). Through the iron line survey observations, Yamauchi et al. discovered a remarkable condensation of the iron line intensity near (≤ 2◦) the Galactic center, as shown also in figure 2.3. They reported that the shape of this region is an ellipse of 1◦.8×1◦.0 tilted by 21◦ with respect to the Galactic plane, and the typical temperature is about 10 keV. The total thermal energy of hot plasma in this region is estimated as 53 4 (4-8)×10 ergs, and the total mass of hot gas (1-2)×10 M⊙ . Possible explanations of this hot gas concentration includes the integrated emission of star forming regions, supernova explosions in the tenuous ambient medium or a single large explosion at the Galactic center. However, definite conclusion is not yet obtained.

Figure 2.3: (a) Longitudinal distribution of the continuum intensity in the 1.1-18.5 keV band obtained with Ginga. It consists of many discrete sources, and the unresolved GRXE. (b) That of the 6.7 keV iron line. (c) An enlarged plot of (b) The solid line represents the best-ﬁt intensity proﬁle (Yamauchi et al. 1993).

9 2.2.2 ASCA results ASCA has been utilized extensively to survey the Galactic center region, and many discrete sources have been resolved from the complex diﬀuse emission. Figure 2.4 shows the very central region of the Galaxy obtained with ASCA (Koyama et al. 1996; Sakano et al. 2000). At the most center, within the 1◦.8×1◦.0 ellipse discovered by Ginga ??, ASCA found a narrower region of 3′×2′ centering on Sgr A∗ where the X-ray surface brightness becomes maximum. The energy spectrum of the diﬀuse emission from the entire ∼1 square degree were characterized by many emission lines from helium-like and hydrogen-like ions of Si, S, Ar, Ca, and Fe, with a continuum described by a thermal bremsstrahlung of ≥ 10 keV (Koyama et al. 1996; Maeda et al. 1998).

Figure 2.4: X-ray image of Figure 2.5: Hard X-ray (3- the Galactic center in 2-10 keV 10 keV) image of the Galac- band obtained with the ASCA tic center region obtained with SIS (Koyama et al. 1996). ASCA (Sakano 2000).

2.2.3 Superposition of Dim Point Sources Against the opinion that the diffuse X-ray emissions in the Galactic center region are truly diffuse emissions, another opinion is supported that the diffuse X-ray emissions can be explained by superposition of a lot of dim X-ray point sources. Since 2000, X-ray imagers with good spacial resolutions of 0.5′′ onboard Chandra satellite have been discovering many X-ray point sources in the Galactic center region (Muno et al. (2006), Figure 2.6, 2.7). The fluxes integrated over all the discovered X-ray point sources can explain only 40 % of whole the detected fluxes. The observational team insists that it is difficult to explain 100 % from X-ray point sources showing the correlations between luminosity and

10 Figure 2.6: X-ray image of the Galactic center region with Chandra.

Figure 2.7: Distribution of X-ray point sources in the Galactic center region detected with Chandra. number density of the discovered sources (Log-N log-S) even if the detection limit becomes better. While, 90 % of the discovered sources are estimated to be Cataclysmic Variables (CVs). The spectra from the CVs are represented by the thermal bremsstrahlung with a temperature of 10 – 40 keV, and therefore can explain the X-ray spectra obtained with ASCA satellite. Hence, assuming that there are a lot of dim point sources whose fluxes are lower than the detection limit, the superposition of point sources possibly can explain both the fluxes and the spectral shapes of the diffuse emission.

11 2.3 Non-thermal Hard Tail Associated with the Milky Way

Another important scientific prospect obtained with Ginga is the discovery of the hard-tail in the GRXE energy spectra as shown in figure 2.8 (Yamasaki et al. 1996a, Yamasaki et al. 1996b, Yamasaki et al. 1997). They discovered that the GRXE spectra obtained with the Ginga LAC cannot be explained adequately with a single-temperature plasma emission, and obtained the best fit with two emission components, a thin thermal plasma of 3.1±1.4 keV and an additional power-law component of photon index ∼1.6±1.1. Since a single power-law extrapolation of the hard-tail connects smoothly to the diffuse Galactic soft γ- ray spectrum (right panel of figure 2.8) which is thought to arise via bremsstrahlung from low-energy cosmic electrons, we presume that some acceleration process is continuously supplying these electrons around the Galactic plane.

Figure 2.8: The hard-tail observed with Ginga (Yamasaki et al. 1997). (Left) The energy spectrum averaged over all scans shown in ﬁgure ??, together with the absorbed optically thin hot plasma model. A discrepancy in the high energy band is clearly shown. (Right) Wide-band spectrum of the Galactic emission from the hard X-ray to soft γ-ray band. A single power-law model of photon index 2.1 is indicated by a dash-dotted line.

12 Figure 2.9: The soft X-ray spectrum of the Galatic center region with Suzaku.

2.4 Suzaku results

We review the following two papers before this thesis based on the analysis of the Suzaku observational data of the Galactic center region. These observations have been performed in the test observational term (within a year from the launch in 2005), and 5 positions of |l| < 1.0◦ have been observed. Koyama et al. (2007c) have performed the precise plasma diagnosis utilizing the observational data (|l| < 0.3◦, |b| < 0.2◦ ) of the X-ray CCD sensors (XIS). They show that the hot plasma is in ionization equilibrium from the energy of the line center and the temperature is 6.5 keV from the line ratio. While, the continuum component of the XIS spectrum is explained by a thermal bremsstrahlung of > 10 keV. Hence, they suggest that there is another component explained by a power-law model, that is the non-thermal emission, from this contradiction (Figure 2.9). However, the spectral shape of the power-law component is not determined at all in only the soft X-ray band. Koyama et al. (2007c) also show that the Galactic diffuse X-ray emission can not be explained by only the superposition of dim point sources from spatial distributions. Figure 2.10 shows the surface brightness distribution of 6.7 keV line compared with the distribution of integrated point sources fluxes. Since these scale lengths are clearly different with each other, they suggest that the diffuse emission can not be explained by the point sources even if dimmer sources are detected with more sensitive imagers in the future. First result based on the Suzaku hard X-ray, Yuasa et al. (2007) have analyzed the observational data with the Hard X-ray Detector (HXD) and discovered the hard X- ray emission which can not be explained by the known 6 hard X-ray point sources. In the estimation of the contributions from the hard X-ray sources, the method described in this thesis is adopted. In addition, we have analyzed the wide band spectra of XIS

13 Figure 2.10: The 6.7 keV line fluxes distribution obtained by Suzaku (squares). And, the integrated point source fluxes in the 4.7 – 8 keV band are plotted by the crosses. and HXD and shown that the power-law component which is suggested by Koyama et al. (2007c) extends to 40 keV. However, there are a lot of hard X-ray point sources in the region of |l| < 1.0◦, and therefore, the systematic errors of the estimations of the contamination from the point sources are large. Hence, the spectral shape of the power-law component can not be determined precisely. Moreover, the hard X-ray surface brightness distribution can not be obtained from only the observational data of |l| < 1.0◦ since the Galactic diffuse X-ray emission is considered to extend to |l| = 2◦ from the 6.7 keV line distribution (Maeda 1998). These results suggest the existence of non-thermal emission in center area of the Galactic center region. The investigation of the hard X-ray emission from the region close to the Galactic center is very important. In this thesis, we extend the analysis of the Suzaku observational data to the position of |l| = 2◦, and aim to obtain the spectral shape and flux of the power-law component without the large systematic errors of the point sources estimations. If we could determine the power-law index, we could constrain the origin of the emission based on the theoretical ground.

14 Chapter 3

The X-Ray Observatory Suzaku

3.1 The Suzaku Spacecraft

The fifth Japanese X-ray astronomy satellite, Suzaku (Mitsuda et al. 2007), was launched on 10 July 2005 with the M-V launch vehicle from Uchinoura Space Center of Japan Aerospace Exploration Agency (JAXA). After the launch, Suzaku first deployed the solar paddles and the extensible optical bench (EOB), and performed ∼ 10 days of the perigee- up orbit maneuver to get into a near circular orbit of 570 km altitude with an inclination of 31◦. The orbital period of the satellite is about 96 minutes. The schematic view and the side view of the Suzaku spacecraft are shown in Figure 3.1 and Figure 3.2, respectively. The five sets of X-ray mirrors are mounted on top of the EOB and five focal plane detectors and a hard X-ray detector are mounted on the base panel of the spacecraft. The spacecraft weight at launch was 1706 kg and its length is 6.5 m along the telescope axis after the deployment of the EOB. The electronics boxes of both the spacecraft bus and the scientific instruments are mounted on the side panels of the spacecraft. The attitude is stabilized by four sets of reaction wheels with one redundancy, while the attitude is measured by three gyroscopes and two star trackers. The spacecraft pointing accuracy is ∼ 0′.24 with a stability better than 0′.022 per 4 sec (a half of typical exposure time of CCD cameras). The pointing direction of the X-ray telescope presently has additional uncertainty and temporal variations due to thermal distortion of the spacecraft structure (Serlemitsos et al. 2007). The normal mode of operations will have the spacecraft pointing in a single direction for at least 1/4 day, which corresponds to a good observation time interval of ∼ 10 ks. With this constraint, most targets will be occulted by the Earth for about one third of each orbit, but some objects near the orbital poles can be observed nearly continuously. Observations are also interrupted by passages of the South Atlantic Anomaly (SAA), in which the particle background drastically increases. The scientific payload of Suzaku initially consisted of three distinct co-aligned detectors. There are four X-ray sensitive imaging CCD cameras, X-ray Imaging Spectrometers (XIS; Koyama et al. 2007a). Three of the XIS sensors are made of front-illuminated

15 Figure 3.1: A schematic picture of the Suzaku satellite in orbit. (Courtesy of ISAS/JAXA)

Figure 3.2: A side view of the Suzaku satellite.

16 (FI; energy range 0.4–12 keV) CCDs and the other is back-illuminated (BI; energy range 0.2–12 keV). Each XIS sensor is located in the foci of a X-ray telescope (XRT; Serlemit- sos et al. 2007) The second instrument is the non-imaging, collimated detector for higher energies, Hard X-ray Detector (HXD; Takahashi et al. 2007). The HXD extends the bandpass of the Suzaku observatory by more than an order of magnitude with its 10– 600 keV bandpass. The last instrument, X-Ray Spectrometer (XRS; Kelley et al. 2007) is the ﬁrst orbiting X-ray microcalorimeter spectrometer. The early veriﬁcation phase of the mission demonstrated that the instrument was working properly and that the cryogen consumption rate was low enough to ensure a mission lifetime exceeding three years. However, the XRS is no longer operational since the liquid He cryogen was completely vaporized two weeks after opening the dewar guard vacuum vent. The XRS and the XRT dedicated to it will not be discussed further in this thesis.

3.2 X-Ray Telescope (XRT)

The Suzaku X-Ray Telescopes (XRTs) are thin-foil-nested Wolter-I type telescopes, which are also utilized in ASCA (Tanaka, Inoue, & Holt 1994), XMM-Newton (Jansen et al. 2001), Swift (Gehrels et al. 2004), and some other missions. These are grazing-incidence reﬂective optics consisting of compactly nested, thin conical elements. The XRT of Suzaku is made of very thin (∼ 178 µm) foils to achieve light weight and high throughput, with moderate imaging capability in the energy range of 0.2–12 keV. Four XRTs onboard Suzaku (XRT-I0 to XRT-I3) are used for the XIS. A photograph of an XRT is shown in Figure 3.3. An XRT is a cylindrical structure, having the following layered components:

1. a thermal shield at the entrance aperture to avoid temperature gradient;

2. a pre-collimator mounted on metal rings for stray light elimination;

3. a primary stage for the ﬁrst X-ray reﬂection;

4. a secondary stage for the second X-ray reﬂection;

5. a base ring for structural integrity and interface with the EOB of the spacecraft.

All these components, except the base rings, are constructed in 90◦ segments, called “quadrants”. Four of the quadrants are coupled together by interconnect-couplers and also by the top and base rings. The telescope housings are made of aluminum for an opti- mal strength to mass ratio. Each reflector consists of a substrate also made of aluminum and an epoxy layer that couples the reflecting gold surface to the substrate. Table 3.1 summarizes the specifications and the characteristics of the XRTs. The angular resolutions of the XRTs are about 2.0′, expressed in terms of half-power diameter (HPD), which is the diameter within which half of the focused X-rays are enclosed. The

17 Figure 3.3: A Suzaku X-ray telescope (XRT). angular resolution does not significantly depend on the energy of the incident X-ray in the energy range of Suzaku, 0.2–12 keV. The effective areas are typically 440 cm2 at 1.5 keV and 250 cm2 at 7 keV per telescope. The focal lengths are 4.75 m. Individual XRT quadrants have their own focal lengths deviated from the design values by a few cm. The optical axes of the quadrants of each XRT are aligned within 2′ from each other. The field of view for XRTs, defined as the full-width-at-half maximum (FWHM), is about 20′ at 1 keV and 14′ at 7 keV. The optical axis of each XRT was determined by observing the Crab Nebula at various off-axis angles. Hereafter all the off-axis angles are expressed in the detector coordinate system (Det-X, Det-Y) (Ishisaki et al. 2007). The result of the determination is shown in Figure 3.4. Since the optical axes moderately scatter around the origin, it was adopted as the XIS-nominal position. On the other hand, the optical axis of the HXD-PIN detector deviates by ∼ 5′ in the negative Det-X direction. Because of this effect, the observation efficiency of the HXD-PIN at the XIS-nominal position is reduced to ∼ 90% of the on-axis value, and another pointing position, HXD-nominal position, is provided for HXD-oriented observations at (Det-X, Det-Y) = (−3′.5, 0′). At the HXD-nominal position, the effciency of the XIS is ∼ 88%. Verification of the imaging capability of the XRTs were made with the observation of a moderately bright point source, SS Cyg. Figure 3.5 shows the image, Point-Spread Function (PSF), and Encircled-energy fraction (EEF) of the XRT modules. The total exposure time used here is 9.1 ks. The obtained HPD is 1′.8, 2′.3, 2′.0, and 2′.0 for XRT-I0, 1, 2, and 3, respectively. It is known that ASCA observations were sometimes hampered by X-rays arriving from sources out of the field of view, which we refer to as stray light. This stray light makes observations of crowed regions like the Galactic center difficult. In order to reduce the effect of the stray light, Suzaku XRT has a pre-collimator in front of the XRT. Observations of stray light were carried out with the Crab Nebula during 2005 August

18 22–September 16 at off-axis angles of (Det-X, Det-Y) = (±20′, 0′), (0′, ±20′), (±50′, 0′), and (0′, ±50′). An example stray-light image is shown in the right panel of Figure 3.6. This image was taken with XIS3 in the 2.5-5.5keV band when the Crab Nebula was offset at (Det-X, Det-Y) = (−20′, 0′). The left and central panels show simulated stray light images without and with the pre-collimator, respectively, of a monochromatic point source of 4.5keV being located at the same off-axis angle. The ghost image seen in the left half of the field of view is due to the stray light. Although the stray light cannot be completely diminished at an off-axis angle of 20′, the center of the field of view is nearly free from stray light.

Table 3.1: Speciﬁcations/Characteristics of the XRTs onboard Suzaku

Focal Lenth 4.75 m Weight/Telescope 19.3 kg Geometrical Area/Telescope 873 cm2 Field of Viewa 17′ at 1.5 keV 13′ at 8 keV Eﬀective Areab 440 cm2 at 1.5 keV 250 cm2 at 8 keV Angular Resolutionb 2′ (HPD)

aDiameter of the area within which the eﬀective area is more than 50% of the on-axis value. bMeasured on the ground

19 Figure 3.4: Locations of the optical axis of each XRT module in the focal plane determined from the observations of the Crab Nebula. The dotted circles are drawn every 30′′ in radius from the XIS-nominal position.

Figure 3.5: Image, Point-Spread Function (PSF), and Encircled-energy fraction (EEF) of the XRT modules in the focal plane. The EEF is normalized to unity at the edge of the CCD chip. With this normalization, the HPD of the XRT-I0 thorough I3 is 1′.8, 2′.3, 2′.0, and 2′.0, respectively.

20 Figure 3.6: Focal plane images formed by stray light. The left and middle panels show simulated images of a monochromatic point-like source of 4.51keV locating at (Det-X, Det-Y) = (−20′, 0′) in the cases of without and with the pre-collimator, respectively. The radial dark lanes are the shades of the alignment bars. The right panel is the in-flight stray image of the Crab Nebula in the 2.5-5.5keV band located at the same off-axis angle. The unit of the color scale of this panel is counts per 16 pixels over the entire exposure time of 8428.8s. The counting rate from the whole image is 0.78 ± 0.01 counts s−1 including background. Note that the intensity of the Crab Nebula measured with XIS3 at the XIS-nominal position is 458 ± 3 counts s−1 in the same 2.5-5.5keV band. All the images are binned with 4 × 4 pixels followed by being smoothed with a Gaussian profile with a sigma of 2 pixels, where the pixel size is 24 µm.

21 3.3 X-Ray Imaging Spectrometers (XIS)

The X-ray Imaging Spectrometers (XISs; Figure 3.7), X-ray sensitive silicon charge- coupled devices (CCDs), are operated in a photon-counding mode, similar to those used in the ASCA SIS, Chandra ACIS, and XMM-Newton EPIC. In general, an X-ray CCD converts an incident X-ray photon into a charge cloud, with the magnitude of charge proportional to the energy of the absorbed X-ray. This charge is then shifted out onto the gate of an output transistor via an application of time-varying electrical potential. Thus, a voltage level (pulse height) proportional to the energy of the X-ray photon is read out. The four Suzaku XISs are designated as XIS0, XIS1, XIS2, and XIS3, located in the focal plane of XRTs; XRT-I0, XRT-I1, XRT-I2, and XRT-I3, respectively. In an XIS camera, there is a single CCD chip with an array of 1024 × 1024 pixels, and covers an 17′.8 × 17′.8 region on the sky. The pixel size is 24 µm × 24 µm, and the the size of the whole chip is 25 mm × 25 mm. One of the XISs, XIS1, uses a back-illuminated (BI) CCD, while the other three use front-illuminated (FI) CCDs. Since the BI CCD has no gate structure on its illuminated side, XIS1 more sensitive to soft X-rays than the other XISs (see Figure 3.8). Table 3.2 is the summary of the specifications and characteristics of XIS. Since the Suzaku launch on 2005 July 10, the XIS has been working properly, and, the on-board performance and the strong points of the XISs have been demonstrated. The features of Suzaku XISs are its low non-X-ray background (NXB) and good energy resolution over a wide spectral band. These features play important roles in observations of diffuse and low surface brightness sources like our Galactic center region. Figure 3.9 shows the NXB spectra of FI and BI sensors when the XIS observes the dark (night) Earth. The data obtained during the passage thorough the SAA and events in the calibration source area (the 2 corners) have been excluded. The lines at 5.9 keV and 6.5 keV are due to scattered X-rays from the calibration sources. Other than these, many lines, Kα of Al, Si, Ni, Kβ of Ni, and Lα,Lβ, and Mα lines of Au are detected. Thanks to the low-Earth orbit of Suzaku, the NXB is fairly low, especially for the FI CCDs. The flux of NXB depends on the cut-off rigidity (COR): the FI fluxes in the 0.4–12 keV band are about 0.2 counts s−1 and 0.1 counts s−1 for the COR of 4–6 GV and 12–14 GV, respectively, while the BI fluxes in the same COR bands are about 0.6 counts s−1 and 0.3 counts s−1, respectively. Thus for the most accurate NXB subtraction, the data should be selected so that the COR distribution of the NXB observation is the same as that in the source observation. Figure 3.10 is the hard X-ray spectrum from our Galactic center region, where the NXB has been subtracted. We see that three Kα lines from Fe I (6.4 keV), Fe XXV (6.7 keV), and Fe XXVI (6.9 keV) are nicely resolved. In addition, Kα lines of Ni I

22 Figure 3.7: A Suzaku XIS sensor. and Ni XXVII, Kβ lines of Fe I, Fe XXV, and Fe XXVI, and even Kγ lines of Fe XXV and Fe XXVI are found above ∼ 7 keV. This demonstrates reliable NXB subtraction and superior energy resolution. These capabilities of the XIS enable us to perform line- resolved imaging studies even of low surface brightness sources. In general, the energy resolution of X-ray CCDs in orbit have got worse due to the radiation damage. In order to recover the degradation of the energy resolution, Suzaku XIS has a capability of a spaced-row charge injection (SCI), and the SCI technique has been applied since 2006 December. The radiation damage makes “traps” in X-ray CCD pixels. These traps obstruct the event transfer, which causes the degradation. If charges are put in front of X-ray events, these charges works as “sacriﬁcial” events. Therefore, they ﬁll the traps and the X-ray events are transferred smoothly, and the energy resolution can be restored. This is the principle of the SCI. Figure 3.11 shows how the SCI works. The spectra show the emission line of the He-like iron from Perseus cluster in the almost same period. The energy resolution is improved by using the SCI technique. The energy resolution of 205 eV (FWHM) without the SCI is recovered to 157 eV (FWHM) due to the SCI. Thanks to the SCI, the energy resolution of the XIS can remain good for about 2.5 years after the launch. One of the major advantages of the XIS over ASCA, Chandra, and XMM is the provision of the SCI.

23 Figure 3.8: Eﬀective area of one XRT + XIS system, for both FI (XIS0, 2, 3) and BI (XIS1) CCDs.

Figure 3.9: The night Earth spectra with the BI and FI CCDs.

Figure 3.10: The X-ray spectrum of the Galactic center. The four XIS data are added. The night Earth background is subtracted.

24 Figure 3.11: Comparison of the He-like Fe Kα line spectrum between with the SCI and without the SCI.

Table 3.2: Speciﬁcations/Characteristics of XIS

Field of View 17′.8 × 17′.8 Energy Range 0.2–12 keV Format 1024 × 1024 pixels Pixel Size 24 µm × 24 µm Energy Resolution ∼ 130 eV (FWHM) at 5.9 keV Eﬀective Areaa 330 cm2 (FI), 370 cm2 (BI) at 1.5 keV 160 cm2 (FI), 110 cm2 (BI) at 8 keV Readout Noise ∼ 2.5 electrons (RMS) Time Resolution 8 s (with Normal Mode)

aOn-axis eﬀective area for one sensor includeing the XRT eﬀective area. The calculations are for a point source integrated over a circular region with a 6 mm (4′.34) radius.

25 3.4 Hard X-Ray Detector (HXD)

3.4.1 Overview The Hard X-ray Detector (HXD; see Figure 3.12) is a non-imaging, collimated hard X- ray instrument sensitive in the ∼ 10 keV to ∼ 600 keV band. The characteristics of the HXD is summarized in Table 3.3. Since the background level sets the sensitivity limit in the hard X-ray band, the HXD is designed to minimize the background by its improved phoswich (acronym for PHOSphor sandWICH) conﬁguration for the energy region above 40 keV and the adoption of newly-developed thick silicon PIN diodes for the energies below 70 keV. Figure 3.13 is the schematic drawing of HXD sensor. Two key techniques are used here: well-type active shield and compound eye conﬁguration.

• Well-type active shield In phoswich counters, two crystals with different decay times are used for the detection part (faster decay time) and the shielding part (slower decay time), and both signals are extracted by a single photomultiplier. The improvement is that the shield is shaped as well, so that it also acts as an active collimator (well-type active shield). This narrows the field of view of the phoswich counter without additional passive material, and results in the main detection part having an active shield of almost 4π of its surrounding (well-type phoswich counter). In the HXD, well-type shield provides very efficient shielding for the the PIN diodes, which are also located at the bottom of the well and are read out independently.

• Compound eye configuration HXD is modular designed, consisting of a number of units. Each well-type phoswich counter unit has a simple shape and operates at a modest count rate by itself. In the HXD, we increase the photon collecting area by placing individual units in a matrix. In this configuration, each unit also becomes an active shield for adjacent units (Compound eye configuration). It is also useful to reduce the possible dead time if parallel processing of each unit could be implemented. For additional shielding for the outer most units, thick anti-coincidence counters are placed surrounding the well units.

The HXD sensor consists of 16 phoswich counters each with 4 silicon PIN diodes, and 20 surrounding anti-coincidence shield counters. The main detection part of phoswich counters is a Gadolinium silicate crystal (GSO; Gd2SiO5(Ce)) buried deep in the bottom of the Well-shape Bismuth germanate crystal (BGO; Bi4Ge3O12) (hereafter we refer it as Well-counter unit). The other main detectors in Well-counter units, PIN diodes, are placed just above the GSO detectors (see Figure 3.15), and read out independently. The eﬀective area of 16 Well-counter units are shown in Figure 3.16. The anti-coincidence

26 Figure 3.12: Hard X-ray Detector (HXD) onboard Suzaku shield counters (Anti-counter units) are also made of BGO. All the 16 Well-counter units and 20 Anti-counter units work independently. Numbering of the Well-counters and Anti-counters, which is frequently referred to in this thesis is shown in Figure 3.14. The Anti-counters also work as an excellent gamma-ray burst monitor with its large effective area for sub-MeV to MeV gamma-rays. The gamma-ray burst location determination is about 5◦. In the HXD, background subtraction is performed by modeling the background spectrum, as was done in the LAC detector in Ginga satellite (Hayashida et al. 1989). This is different from the other hard X-ray instruments such as CGRO OSSE, BeppoSAX PDS (Frontera et al. 1998), and RXTE HEXTE (Rothschild et al. 1998), which incorporates time sliced on-off observations to subtract background. Thus the sensitivity depends on the accuracy of the background modeling. Thanks to the low background level of the HXD, the systematic error of the modeling has less effect on the spectrum of sources than other experiments. The method of background estimation will be discussed in §6.2.

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28 Configuration of Sensor Units (Top View)

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P0 P1 T31 W33 W32 W23 W22 T14 P3 P2 T30 T24 T23 T22 T21 T20 PIN in Well 1 unit

Figure 3.14: Numbering of the Well and Anti counter units when HXD Sensor is viewed from the top. There are 16 Well-counter units from W00 to W33 and 20 Anti-counter units from T00 to T34. The Y-direction corresponds to the direction toward Sun when the HXD is mounted to the spacecraft.

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Figure 3.15: A schematic drawing of the Well-counter unit. The silicon PIN diode with dimensions of 21.5 mm × 21.5 mm × 2 mm is placed in the deep well just above the GSO scintillator. The 24 mm × 24 mm GSO scintillator with a thickness of 5 mm is glued to the bottom part of the well-type BGO shield.

29 Figure 3.16: Total eﬀective area of the HXD detectors, PIN and GSO, as a function of energy. Photon absorption by materials in front of the device is taken into account.

Table 3.3: Speciﬁcations/Characteristics of HXD

Field of View 34′ × 34′ (. 100 keV) 4◦.5 × 4◦.5 (& 100 keV) Energy Range 10–600 keV – PIN 10–70 keV – GSO 40–600 keV Energy Resolution ∼ – PIN 4.0√ keV (FWHM) – GSO 7.0/ EMeV % (FWHM) Eﬀective Area ∼ 160 cm2 at 20 keV ∼ 260 cm2 at 100 keV Time Resolution 61 µs or 31 µs

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3.4.2 HXD-PIN Detectors

The silicon PIN diodes cover the lower energy region of the HXD bandpass, from 10 keV to 70 keV. For the HXD data, only the data from PIN detectors are used in this thesis. Figure 3.17 shows the photograph and the schematic picture of the PIN diode. The geometrical area of the PIN diode, including the guard ring structures, is 21.5 mm × 21.5 mm with a thickness of 2 mm. The leakage current is less than 2.2 nA at the operation condi- tions (500 V and −20 ◦C). In order to minimize the background gamma-ray contamination for the PIN, we have selected low-background material for cables and package. The adoption of ceramic with a purity higher than 96% for the PIN package was required to avoid continuum gamma-ray background due to β-decay electrons from potassium in the material. Figure 3.18 demonstrates the superior performance of the PIN diode. This is a spectrum of X-rays/gamma-rays from 241Am obtained with a PIN diode on the ground. As shown in this figure, an energy resolution of 2.8 keV was achieved and the energy threshold was ∼ 10 keV. The energy resolution of the PIN preamplifier is 1.0 keV in a load-free condition; this degrades to 1.6 keV by the capacitive noise, and another 1.0 keV is due to the leakage current noise. The remaining ∼ 2.1 keV is possibly caused by electronic noise. Field of view (FOV) of the HXD-PIN is collimated with fine collimators in order to reduce contamination of Cosmic X-ray Background (CXB) and point sources which are not targets. Having a narrow field of view is the most effective method to reduce the background contamination. For this purpose, passive shields called ”fine collimators” are inserted to the BGO well-type collimator above the PIN diodes. The fine collimator is made of 50 µm thick phosphor bronze sheets arranged to form a square array of 8 × 8 channels each of 3 mm width and 300 mm length. The fine collimators confine the FOV of PIN diodes even narrower than that restricted by the active shield of BGO. The FOV defined by the fine collimators is 34.2′ × 34.2′ (FWHM) square in the case that these are produced according to the designed precisely. A cross section of the FOV is an ideal

31 Figure 3.18: A PIN spectrum of 241Am obtained on the ground. The temperature is −20 ◦C triangle with FWHM of 34.2′. Calibrations of the angular response is described precisely in the next chapter.

3.4.3 In-Orbit Calibration Calibrations of gains and energy scale and measurements of energy resolution are essential in constructing the instrumental response function, which in turn is needed to reconstruct incident spectra from the observed pulse-height distributions. Before the launch, the energy scales of the 64 PINs were precisely measured using the standard gamma-ray sources, within ∼1% accuracy (Takahashi et al. 2007). These energy scales are not expected to change significantly after the launch, because neither the charge collection efficiency of the PIN diodes nor the capacitance of the charge sensitive amplifiers is sensitive to the environmental changes. Nevertheless, the energy scale is so important that it should be accurately reconfirmed using the actual data. In order to determine the energy scale, line features in the spectrum are useful. Since the line emissions can be rarely observed from celestial objects in the bandpass of HXD- PIN, X-ray line emissions generated within the detector are utilized. Figure 3.19 illus- trates the events used for the calibration. When an X-ray photon is absorbed in GSO, a

ﬂuorescent X-ray photon from gadlinium (Gd-Kα: 42.7 keV) sometimes escape from the GSO and is absorbed in the PIN detector with some probability. Figure 3.20 shows the positions of Gd-K peak channels of 64 PINs between two periods which are 1-2 and 7-9 months after the launch. It is conﬁrmed that the gains of HXD-PIN have stayed constant within one ADC channel which corresponds to ∼ 0.4 keV (Tanaka 2007).

32 X-rays X-rays

Absorption

Bi-K Gd-K PIN PIN

GSO Absorption GSO BGO Well BGO Well

Figure 3.19: Illustration of an event in which Gd-K X-ray is detected with a PIN diode (left) and in which Bi-K is detected.

Figure 3.20: Comparison of the Gd-K peak channels of the 64 PINs between two periods which are 1–2 and 7–9 months after the launch (Tanaka 2007).

33 Chapter 4

Angular Response of HXD-PIN and A New Method for Flux Estimation

Calibrations of the angular response is essential in estimating fluxes of point sources in the FOV of HXD-PIN, especially in care of crowded region like the Galactic center. In this chapter, first, we describe in-orbit calibrations of the angular response. Then, we introduce a method to estimate the flux of point sources in the FOV by use of differences of individual optical axes of 64 fine-collimetors of the PIN detector.

4.1 Angular Response of HXD-PIN

4.1.1 Fine-Collimator

In the energy range of the PIN diode, the dominant background component is the cosmic X-ray background. Therefore, having a narrow FOV is the most effective method to reduce the background contamination. For this purpose, passive shields called “fine collimators” are inserted in the BGO well-type collimator above the PIN diodes. The fine collimator is made of 50 µm thick phosphor bronze sheets arranged to form a square array of 8 × 8 square channels each of 3 mm width and 300 mm length. Both the BGO collimator and the fine collimator define the FOV of the Well-counter unit. Because of the finite thickness, the FOV changes with the photon energy. Below ∼ 100 keV, the fine collimators confine the FOV of PIN diodes much narrower than the BGO collimator and define a 34′ × 34′ Full-Width Half-Maximum (FWHM) square FOV. Above ∼ 100 keV, the fine collimators become transparent and the BGO active collimator defines a 4.5o× 4.5 o FWHM square opening. The angular response of HXD-PIN is determined by the fine-collimator . Functional form the angular transmission of the fine collimator is plotted in Figure 4.1 with respect to the incident angle of X-rays, and we call this as an angular response of HXD-PIN. The function is determined to be almost a triangle shape with a width of 34.2′ (FWHM, Full Width at Half Maximum), which is calculated from the width of 3 mm and length

34 of 300 mm of the opening part of the fine collimator. Hereafter, we define a direction at which the transmission becomes maximum as the ”light axis” of the fine-collimator.

4.1.2 Angular Response

Angular response, the transmission efficiency of the collimator, is one of most important issues of calibration. Without accurate number, we would not be able to derive the flux from sources. The angular response of the fine collimator depends on the direction of its light axis with respect to the pointing direction of the satellite. If the alignment of the collimator is different with each other, we have to take these differences into account when we calculate the angular response of the PIN detector. Meanwhile, the angular response dose not depend on the energy range of HXD-PIN (below 80 keV) (Figure 4.1), therefore we do not have to include the energy dependency in this thesis. In the actual observational data, the angular response is characterized by the equation;

F [i] R[i] = , (4.1) f

where R is the angular response for each fine collimator, i is ID of the PIN corresponding to the collimator, F is the flux arrived at the surface of the PIN detector and f is the expected flux when we locate the target in the direction of light axis. The functional form of the angular response with respect to the pointing direction of the satellite is defined by the shape of the fine collimator as schematically shown in Figure4.2. In the figure, an effective area is indicated as ”not-shadowed” region, and therefore the ratio of observed flux and the real flux corresponds to the ratio of the area of ”not-shadowed” region divided by the total effective area of the PIN detector. Once we obtain an offset of the light axes of the fine-collimators from the pointing direction of the satellite, we are able to calculate the angular response for a source from any directions in the FOV by using a triangle-shaped function defined by the mechanical structure and its thickness. Figure 4.3 shows the 1-D profile of the angular response for the position of the point source in the FOV. And, Figure 4.4 shows the 2-D one.

4.2 In-Orbit Calibration of the Angular Response

4.2.1 Calibration of Light Axes

Although the directions of 64 ﬁne collimators have been ﬁnely adjusted before launch, there still remain slight (∼ 1′) variations between individual collimators. Therefore, we tried to obtain more accurate directions of the individual light axes by performing a series of in-orbit calibration observations.

35 Figure 4.1: Angular transmission function of the ﬁne collimator calculated at azimuth angle of 0◦. The 0◦ azimuth angle is deﬁned along the positive X-axis in Figure 3.14.

36 Target

Light axis

1 pixel / 8 x 8 pixels Side view of ﬁne collimator (X-axis)

Shadow of X-axis (Bx)

Bottom view A Shadow Shadow of Y-axis (By)

Figure 4.2: Schematic representation of the angular response of a PIN. Inci- dent X-rays reach at the non-shadowed region.

37 ARF

1.0

ARF[i]

34.2’

Target Light axis of the PIN Nominal position of the satellinte

Figure 4.3: Schematic perspective of determination of the ratio.

40 1

30 0.8 20

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Figure 4.4: Schematic perspective of an eﬀectve area of a PIN in an observation in which a direction to a target is diﬀerent from that of a light axis.

38 Figure 4.5: Distribution of light axes of the individual ﬁne-collimators measured in the ground test.

Before the launch, the light axes of the 64 fine collimators have been measured by means of optical laser light and radio-isotope sources. Figure 4.5 shows their angular offsets measured by scanning 31 keV gamma-ray line from 133Ba. We define the origin (δx = 0, δy = 0) at the mean of the offsets of each light axes of 64 fine-collimators. According to the ground calibration, the collimators are aligned up with an accuracy of 3.5′ (FWHM). This ensures an effective transparency of 90%, when a target is placed at the mean direction of optical axes of the 64 fine-collimators. Distribution of the light axes, shown in Figure 4.5, leads to deviation of effective area for each PIN detector with a range of ∼ 20 % . When we sum up the flux measured by each PIN detector, we have to apply these factors to the flux obtained from individual detectors. In the ground tests, it was difficult to prepare a bright and parallel hard X-ray beam which is necessary to calibrate the light axes precisely. In addition, the light axes may have changed by the vibrations of the launch and thermal stress. Therefore, we have to perform in-orbit calibration of the light axes. We used Crab Nebula as a calibration target, since the source is bright in the hard X-ray band from 10 keV up to 50 keV and fluxes of the source are known to be stable in time. To obtain the actual shape of the transmission function used to reconstruct the angular response for each collimator, we have performed multiple observations of the Crab with several offset angles separated with ∼ 10′. Table 4.1 shows the observation log of the Crab-scan observations. In order to determine the light axes of the individual fine-collimetors, we have to know the peak positions of the triangle along both the X-axis and the Y-axis. Therefore, we performed the scanning observation along both axes with offset angles of 0′, ±3.5′, ±7.0′, ±10.0′and ± 20.0′.

39 We determined these pitches such that we could obtain the peak position with a suﬃcient accuracy compared to the size of the triangle of 34.2′ (FWHM). Counting-rate distribution is shown in Figure 4.6. with an energy range of 15–40 keV. This energy rage is chosen to reduce systematic uncertainties to the angular response, due to possible contamination of thermal noise around 10 keV. We apply the standard screening criteria (ver. 2.0) which is described in chapter 6.

Table 4.1: Summary of the Crab observations of various angles.

Sequence Angle for the Craba Date Expb number (X, Y) 100007020 (−10′, 0′) 2005/08/22 2.8 100007030 (+10′, 0′) 2005/08/22 2.0 100007040 (0′, + 10′) 2005/08/22 3.9 100007050 (0′, − 10′) 2005/08/22 3.0 100010010 (−3′.5, 0′) 2005/08/24 3.2 100010020 (+3′.5, 0′) 2005/08/25 3.4 100010030 (−20′, 0′) 2005/08/25 7.3 100010040 (+20′, 0′) 2005/08/25 7.8 100010050 (0′, + 20′) 2005/08/25 6.4 100010060 (0′, + 3′.5) 2005/08/25 4.7 100010070 (0′, − 3′.5) 2005/08/26 3.8 100010080 (−7′, 0′) 2005/08/26 3.3 100010100 (0′, − 20′) 2005/08/26 7.2 100010120 (+7′, 0′) 2005/08/26 4.8 100010130 (0′, + 7′) 2005/08/27 4.0 100010140 (0′, − 7′) 2005/08/27 3.3 100023020 (0′, 0′) 2005/09/13 12.1

aThe oﬀset angles for the Crab in the satellite nominal coordinate. bEﬀective exposure of HXD-PIN data.

We determine individual peak positions by ﬁtting the distribution of the counting rate along X-axis and Y-axis with the triangle of a width of 34.2′ (FWHM). Since the actual shape of the angular response can not be modeled by a triangle function around the top, we perform ﬁtting with the data excluding points within 3.5′ from the peak position as shown in the Figure 4.6. We then perform the same procedure for 64 PINs along both the X and Y-axis, and obtain 64 sets of peak positions (X, Y) in total. Two-dimentional distribution of the determined light axes is shown in Figure 4.7.

40 Figure 4.6: An example of the count rates distribution extracted from the scanning observations of the Crab. The PIN ID is W00P3.

Figure 4.7: A distribution of light axes of the individual PINs’.

41 The “light axis of HXD-PIN”, namely, the average direction of 64 light axes, at which the total transparency of fine-collimetors become maximum, is determined as (X, Y) = (−4.03′, −0.62′) and root mean squares of the distribution are 1.58′ for the X-axis and 1.70′ for the Y-axis. The difference of light axes between HXD-PIN and XIS exists, since detected photons with HXD-PIN decreases by 13% in the case that the nominal axis of the satellite is directed at the target source. Therefore Suzaku team prepared the second nominal position of the satellite ”HXD-nominal” with the position of (X, Y) = (−3.5′, −0.0′) in addition to the first that of ”XIS-nominal” with (0.0′, 0.0′).

4.2.2 Fine Tuning of the Angular Response Utilizing the measured directions of the light axes for all fine-collimetors, we reconstruct an actual shape of the angular response as shown in Figure 4.8. The figure is created by plotting 64 points from each the Crab scanning observation. The energy range is 16–40 keV. The sky coordinate positions of the individual data points correspond to directions of the light axes of the individual fine-collimators in the actual observations. The upper panel of the figure is the plot drown along RA-axis of sky coordinate and the lower panel DEC-axis. The plots of the RA and DEC-axis are created with the observational data of the Crab scanning observations along Y and X-axis respectively. The vertical axis is detected counting rates during the observation the individual PINs. We note that the correspondence of X-axis and Y-axis with axes of RA and DEC is inverted since the roll angles of these observation are ∼ 90◦. The center of the triangle certainly agree with the position of the Crab. In Figure 4.8, the counting rates are corrected with difference of relative effective areas in the energy range of 16−40 keV. The effective area of each PIN diode is calculated from a thickness of a depletion layer of the detector. The PIN diodes of HXD-PIN are as thick as ∼ 2.0 mm and a bias voltage around 700 V is required to obtain full depletion of Si (Ota et al. 1999). At the nominal operation voltage of ∼ 500 V, the actual thickness of the depletion layer could vary among the 64 PINs. The thicknesses of the individual PIN diodes were tuned using an observational data of the Crab (Tanaka 2007). And then, the difference of the thicknesses should be taken into account when the energy response of the individual PINs. Using the individual energy responses (ver. 2007-09-14), we extracted the relative effective areas as shown Figure 4.9. The relative effective areas were normalized to 1.0 when the average of 64 PINs was set to 1.0. The counting rates of Figure 4.8 were calculated with dividing the raw counting rates by the relative affective areas. As shown in Figure 4.8, it is clear that the light axes of individual PINs have been calibrated adequately and the function of angular response is close to the triangle. However, we confirm that the top of the triangle has a round shape in the actual data. In addition, the width of the derived function is narrower than the ideal triangle and becomes 33.0’ (FWHM). These differences can be explained as an effect caused by small distortions in

42 each fine collimator. The round shape near the top of the triangle could be explained by small differences of directions, ∼ 1′, among 8×8 pixels in ”one” fine-collimator. The narrower width could be explained by surface roughness and distortion (∼ 50 µm) of the phosphor bronze used in the fine-collimator. We use the function extracted from the actual data as the angular response for point sources in this thesis, as shown in Figure 4.10. The map of the corrected angular response summed with 64 PINs is also shown in Figure 4.11.

4.3 New Method for Flux Estimation

Distribution of corrected flux at different offset angles, shown in Figure 4.8 , suggests that the accurate estimation of a flux of a point source in the FOV is possible by using the distribution of the count rate from individual PIN detectors, if the position of the sources is given. The flux, the height of the triangle, can be calculated by solving a simple linear equation which consists of position of the peak and leaning of the data points. The calculation is afford to say as a simple imaging with a large PSF of 33′ and small FOV of several squares of minutes. We have developed a method to estimate the flux of the source even if the FOV is filled with another uniform emission such as the diffuse emission of the Galactic center region. As a summary, the flux estimation method introduced here is to solve the linear equations explained above for the individual PINs. When the FOV is filled with an uniform emission, an equation is provided as follows;

F [i] = R[i]f + B, (4.2) where F is a flux arrived at the surface of the PIN detector, R is angular response, i is PIN’s ID, f is expected flux in the case that the target exists in the direction of light axis and B is an uniform background emission. In this equation, known parameters are F [i] and R[i], while unknown parameters f and B. Since the number of the unknown parameters is two, we are able to obtain values of the variables as long as two individual equations of equation 4.2 from the observational results of two PINs’ are given. Since statistics of the two PIN’s fluxes are usually not enough to calculate f of the target, we utilized equations obtained from summed data of 64 PINs. In order to obtain the most effective statics, we define an equation

∑near ∑far ∑near ∑far F [i] − F [i] = R[i]f − R[i]f, (4.3) i i i i where ”near” and ”far” are groups of PINs whose light axes are nearer and farer from a direction of the target. The ﬂux of the target is given as below equation, ∑near ∑far ∑near ∑far f = ( F [i] − F [i]) / ( R[i] − R[i]). (4.4) i i i i

43 It is assumed that the uniform background emission is constant for the individual PINs. An example of a correlation between R and counting rates of the individual PINs are shown in Figure 4.12. The figure was created by using the actual data of one of the Crab scanning observations. The sequence number of the observation is ”10007020” whose position to the Crab is (X,Y ) = (+3.5’, 0′). The position is counter side of the light axis of HXD-PIN for the Crab. It is expected that we are able to clearly confirm the correlation between the ratio of the angular response and the actual counting rates using the data of the position, in which the lean of the function is constant and the statics of the counting rates is good. In the figure, the horizontal axis means the ratio of the angular response of the individual PINs to the targets in the FOV. Meanwhile, the vertical axis means the counting rates of the individual PINs corrected with the relative effective areas. An energy range of the counting rates is 16−40 keV. We confirmed that there are clear correlations between the ratio and the counting rates. In order to confirm if the method can be utilized realistically, we calculated estimation limits of this method. An uncertainty of the calculated flux is determined by statistics of the NXB and fluxes from the source. And then, the detection limits are defined by the significance for the statics of the NXB. At the observation of ”10007020”, an expected counting rate of the NXB is 0.26 counts/sec in an energy range of 16–40 keV, for example. From the rate, counting rate of each PIN is calculated to 4 ×10−3 counts/sec/PIN roughly. We assumed this rate as the typical counting rates and calculated the limits of the flux estimation method. When an exposure of an observation is EXP sec, an integration of count of the NXB is equal to 4 ×10−3 × EXP counts/PIN. At the equation 4.4, the flux of the source is calculated by the counting rate F from the target and the ratio of the angular response R. In the actual analysis, F is calculated by subtracting counting rate of the NXB from the detected counting rate. Therefore, statics of F are determined by the statics of the NXB and counting rate∑ form the target. And√ then, the statical uncertainty√ of the component in × × −3 × × the equation, i F [i], is calculated by 32 (4 10 EXP ) = 0.13 EXP since √the number of the PIN is 32 for each group. And then, changed to a unit of counts/sec, 0.13/EXP . Finally, the uncertainty of the estimated flux, ∆f, is calculated as follows,

√ ∑near ∑far ∆f = 2 × 0.13/EXP / ( R[i] − R[i]). (4.5) i i

The ratio R is calculated for each position, and a map of ∆f calculated for each meshed position is shown in Figure 4.13. We created the figure assuming that the exposure is 40 ksec. The units of the map is ”mCrab” which is corresponds to 4.6×10−4counts/sec in the energy range of 16–40 keV. In an observation of exposure of 40 ksec, the uncertainty at the most sensitive position is ∼ 1 mCrab at 1 σ significance. In this figure, the sensitivity in∑ the center of∑ FOV is worse since the difference of counting rates between near − far two groups,( i F [i] i F [i]), is smaller because of the rounding top of the angular

44 Figure 4.8: Reconstructed shape of the angular response with the Crab scanning observations. response While, the sensitivity around the end is worse since the transmittance is small and then the static of the detected counting rate is not good. Utilizing this method, we estimated the flux of the Crab using the results of the scanning observations listed in Table 4.1. Figure 4.14 shows the estimated fluxes of the Crab compared among the scanning observations for X and Y-axis. We confirmed that the fluxes are estimated adequately within statistical uncertainty. Since the direction of light axis of HXD-PIN is (X,Y) = (−4.03′, −0.06′) and the rounding top of the angular response is positioned at the direction, the uncertainty of the observation in (X,Y) = (−3.5′, 0′) is worse. Performing the flux estimation method every energy bin, we are able to obtain spectrum of the target source. Figure 4.15 shows spectra of the Crab extracted with the estimation method and regular means from the observational data of ”100023020” whose nominal position of the satellite is XIS-nominal of (X, Y) = (0′, 0′). The regular means is described in chapter 6. An energy response of the spectrum with the regular means is summed for 64 PINs. Meanwhile spectrum with the estimation method is defined as that of a PIN detector. Therefore we scaled the spectrum with the estimation method by 64 times. We have confirmed that the two spectra is consistent with each other and we are able to obtain even a energy spectrum with this method if the target is bright enough for a static.

45 Figure 4.9: Distribution of the relative eﬀective area among PINs’.

Figure 4.10: Cross sections of the angular responses. The line is an actual shape and the dashed line is an ideal shape.

46 40 1

30 0.8 20

10 0.6 (’)

Y - 0 T E

D 0.4 -10

-20 0.2 -30

-40 0 -40 -30 -20 -10 0 10 20 30 40 DET-X (’)

Figure 4.11: Actual angular response of the HXD-PIN which is summed up responses of 64 PINs’.

Far Near

Figure 4.12: Correlation between ARF and count rate. (100007020)

47 mCrab 10 40

20 7.5

10 (’)

- 0 5 T E

D -10

-20 2.5 -30

-40 0 -40 -30 -20 -10 0 10 20 30 40 DET-X (’)

Figure 4.13: Detection limit of the simple imaging method with an exposure of 40 ksec.

Figure 4.14: Fluxes of the Crab estimated with the imaging method.

48 iue41:Setu fteCa siae ihteiaigmethod. imaging the with estimated Crab the of Spectrum 4.15: Figure

ratio normalized counts/sec/keV 0.5 1 1.5 2 0.1 0.2 0.5 1 2 5 2 0 c h a n 49 n e l

e n e r g y

( k e V ) Chapter 5

Suzaku Observations

5.1 Overview

Since the launch of Suzaku in 2005, we have performed thirty five observations the Galac- tic Center region (|l| < 2◦.0, |b| < 0.5◦). Most of the observations were performed in 4 terms; from 2005-09-23 to 2005-10-12, from 2006-02-20 to 2006-03-01, from 2006-09-09 to 2006-10-12 and from 2007-03-03 to 2007-03-18. The first two terms were done in SWG (Science Working Group) phase and the last two were done in AO-1 (1st Announcement of Oppotunity) phase. Total net exposure is amount be ∼1 Msec . The details of the observations, exposures, dates and positions of the observations, is listed in table 5.1. The position of each pointing is show in figure 5.1 as an exposure map of the XIS.

5.2 Strategy of the observations

According to results from previous observations, we now have a confidence on the diffuse emission from hot plasma in the Galactic Center region (see Chapter 2). In order to further study the nature of the hot plasma associated with the Galactic Center region and its origin, hard X-ray observations give us crucial information. Furthermore, if the hard X-ray emission can not be explained as the extension of the thermal spectrum of hot plasma, this would indicate the existence of non-thermal emission in the Center of our Galaxy. The temperature of the hot plasma can be determined by either line profiles or the shape of the continuum. The measured temperature from the Galactic Center Region is ∼7 keV, if we use results from the ASCA observation. Therefore, if non-thermal emission would co-exist, it should exceed the thermal emission above ∼10 keV. The combination of XIS and HXD onboard Suzaku gives us the best opportunity to investigate the hard X-ray emission and to study whether it is non-thermal origin or not, through its capability on wide band spectrum coverage, particularly a spectrum from 5 keV up to 50 keV. One of main motivation of the observations is to discriminate the non-thermal emission from the diffuse emission from the Galactic Center Region, for the

50 ﬁrst time. The observation of the Glactic Center Region was originally proposed as a big project for the entire Suzaku team, and there are several objectives in its 1 Ms observation. Here we summarize four strategies related to the investigation of the non-thermal component in the wide band specra.

5.2.1 Determination of the temperature of the hot plasma In order to study an excess of the non-thermal emission, we need to constrain the shape of spectra of the thermal emission. The first step is to determine the temperature of the plasma and its distribution. The temperature is given by an intensity ratio between 6.7 keV and 6.9 keV lines, which corresponds to K-α lines from He-like and H-like Fe ion in optically thin hot plasma. Based on the previous observations of the Galactic Center Region with Ginga, ASCA, Chandra and Newton, we define the region of interest to be |l| < 2◦.0, |b| < 0.5◦. By taking the size of the FOV of the XIS into account, we divide the region into two. Also, in order to increase the robustness of the temperature determined from the two Fe line emissions, we observed the same position twice for some cases. There are four observations in September 2005. One of the two positions was observed again as a calibration of the XIS in March 2006. In total, we have five observations, in which data is mostly used to determine the properties of thermal spectrum which characterize the Galactic Center Region. Exposures of the individual observations were decided to be large enough to resolve the two Fe line emissions clearly and obtain a detailed temperature of the hot plasma. Roll angles during each observation were optimized to the requirement came from by the sun angle.

5.2.2 Monitoring bright transient hard X-ray sources In the Galactic Center region, there are many transient X-ray and hard X-ray sources. Although the collimated instrument, like HXD, has high sensitivity especially for the diffuse emission, it would be difficult to extract fluxes from point sources in the FOV. The smallest FOV of HXD-PIN allows us to reliably estimate the contamination of hard X-ray point sources, since expected number of sources in the FOV is small for most of cases and we have simultaneous data from XIS. The XIS data can be used to constrain the source acticity. Most of known sources in the Galactic Center Region has changes the flux with time scale of a week, therefore, the aim is satisfied with these sources if we could monitor the source activity with XIS during this time scale. The second strategy is, therefore, to monitor especially bright hard x-ray sources with the XIS. We selected three hard x-ray sources with fluxes of above 10 mCrab, which were ”1A 1742-294”, ”KS 1741-294”, ”1E 1743.1-2843”. Positions of the observations were determined to be center of the sources. The observations were performed for each position. The sources are so bright that small exposure of ∼ 3 ksec is sufficient to constrain their

51 spectra and ﬂux.

5.2.3 Observations of molecular clouds In the Galactic Center region, there are several giant molecular clouds which are known to be a strong emitter of neutral Fe line of 6.4 keV. X-ray emission from the clouds itself is an important topic to study the activities of the Galactic Center, but these are out of scope of this thesis. In the Galactic Center Region observation, signiﬁcant amount of exposure time was given in the observations of the cloud, ”Sgr B2”, ”Sgr C” and ”Sgr D”. For the HXD, bright hard x-ray sources are relatively small in these regions, we use the data set from the observations of molecular clouds to study the distribution of hard X-ray emission in both positive and negative side of galactic latidude.

5.2.4 Mapping observations In order to study the origin of hot plasma and also to investigate the hard X-ray emission associated with the Galactic Center Region, it is important to cover the entire region continuously. According to the spacial distribution of 6.7 keV line emissions which indicate an existance of the hot plasma, a scale height of the Galactic diffuse emission is predicted to be ∼ 0◦.5 from a We, therefore, adopted ∼ 0◦.2 for a pitch of positions of observations which is smaller than the scale height if we take the FOV of the XIS into account. We gave a priority to cover to move the FOV along the Galactic longitude since the 6.7 keV line emissions are distributed wider along the Galactic plane than that along the Galactic bulge. These mapping observations were defined to the 4th starategy. There are some other observations in the Galactic Center region, which were proposed under different motivation, but we included those data in this thesis if the data can be accessed.

5.3 Status of the Observations

In all the observations of the Galactic center region used in this thesis, the XIS was operated in the normal clocking mode with no window option. The Space-row Charge Injection mode (SCI) is used in the observations after September 2006 after the mode became available because the deterioration of the energy resolution of X-ray CCDs is recovered if we use the SCI mode. Since one of four sensors of the XIS, ”XIS-2” was lost on 2006-11-09, we do not use the data from XIS-2 for ten observations in March 2007.

52 Table 5.1. Observations of the Galactic Center region.

Step Sequence Positiona Start Expb SCIc 400V-HVd number l b (UTC) (ks) on/oﬀ unit

1 100027010 0.057 -0.074 2005-09-23 07:07:00 34.9 off – 100027020 -0.247 -0.046 2005-09-24 14:16:00 33.4 off – 100037010 -0.247 -0.046 2005-09-29 04:25:00 36.7 off – 100037040 0.057 -0.074 2005-09-30 07:41:00 33.0 off – 501046010 -0.167 0.333 2007-03-10 14:43:00 22.9 on W0,W1 2 100027030 -0.441 -0.389 2005-09-24 11:05:00 1.7 off – 100027040 -0.446 -0.067 2005-09-24 12:40:00 1.6 off – 100027050 0.328 0.010 2005-09-25 17:27:00 1.7 off – 100037020 -0.441 -0.389 2005-09-30 04:29:00 2.1 off – 100037030 -0.446 -0.067 2005-09-30 06:05:00 2.5 off – 100037050 0.328 0.010 2005-10-01 06:21:00 2.0 off – 3 100037060 0.637 -0.095 2005-10-10 12:07:00 65.9 off – 100037070 1.000 -0.100 2005-10-12 07:05:00 8.6 off – 500018010 -0.569 -0.093 2006-02-20 12:30:00 43.8 off – 500019010 -1.091 -0.041 2006-02-23 10:50:00 11.0 off – 501039010 0.780 -0.160 2007-03-03 12:05:00 84.6 on W0,W1 501040010 0.607 0.072 2006-09-21 17:21:00 49.8 on W0 501058010 1.300 0.200 2007-03-14 05:00:00 46.8 on W0,W1 501059010 1.167 0.000 2007-03-15 18:55:00 49.6 on W0,W1 501060010 1.500 0.000 2007-03-17 05:06:00 50.2 on W0,W1 4 500005010 0.428 -0.117 2006-03-27 22:40:00 60.4 off – 100048010 0.057 -0.074 2006-09-08 02:11:00 56.1 off W0 501008010 -0.154 -0.191 2006-09-26 14:14:00 102.6 on W0 501009010 -0.074 0.178 2006-09-29 21:25:00 43.6 on W0 501049010 -1.167 0.333 2006-10-08 10:19:00 16.2 on W0,W1 501050010 -0.833 0.000 2006-10-09 02:19:00 17.5 on W0,W1 501051010 -1.167 -0.000 2006-10-09 13:39:00 19.4 on W0,W1 501052010 -1.5 0.000 2006-10-10 06:44:00 14.9 on W0,W1 501053010 -1.833 -0.000 2006-10-10 21:18:00 18.5 on W0,W1 501057010 -1.167 -0.333 2006-10-11 10:06:00 17.5 on W0,W1 501047010 -0.5 0.333 2007-03-11 03:55:00 17.8 on W0,W1 501048010 -0.833 0.333 2007-03-11 19:04:00 22.3 on W0,W1

53 Table 5.1—Continued

Step Sequence Positiona Start Expb SCIc 400V-HVd number l b (UTC) (ks) on/oﬀ unit

4 501054010 -0.167 -0.333 2007-03-12 08:09:00 21.5 on W0,W1 501055010 -0.5 -0.333 2007-03-12 23:58:00 19.9 on W0,W1 501056010 -0.833 -0.333 2007-03-13 15:40:00 23.1 on W0,W1

aFOV (XIS-nominal) center position.

bEﬀective exposure of the HXD-PIN data.

cXIS’s SCI mode. dHXD-PIN’s units whose bias voltage was 400 V.

. 501046010 501047010 501048010 501049010

501058010 0

50 1009010

. 501040010 0 100027040 501060010 501059010 100027050 501050010 501051010 501052010 501053010 100037050 100027020 1 00037030 500019010 2.000 1.500 1.000 0.500 0.000 100037010 35 590.500018010 359.000 358.500 358.000 100037070 10003706 05 00005010 100027010 501039010 100037040

100048010 0 0 501008010

0 - 501054010 501055010 501056010 501057010

0 100027030

0 100037020

0 500 1000 1500 2000 2500

Figure 5.1: The exposure map of the XIS in the Galactic Center region.

54 Chapter 6

Suzaku Data Analysis and Results

In this chapter, we describe data reduction procedures and results on the analysis. Firstly, we analyze the data obtained with XIS and HXD-PIN independently. Secondly, the contaminations from the hard X-ray sources are subtracted from the HXD-PIN spectra. Thirdly, we perform the joint analysis of the XIS and HXD-PIN spectra to investigate the non-thermal emission from the Galactic center region.

6.1 XIS Data Analysis

6.1.1 Data Reduction

We have used Suzaku data sets processed by the Suzaku data processing pipeline version 2.0. together with the calibration data set of version 2007-08-04. We have screened the data set with the standard event selections as follow. We have filtered out data obtained during passages through the South Atlantic Anomaly (SAA) since Non X-ray Background (NXB) of the XIS was large enough to saturate the sensor. And the data has been filtered with elevation angle to the Earth’s limb below 5◦, or with elevation angle to the bright Earth’s limb below 20◦, in order to avoid the possible contamination of emission of the scattered solar X-rays from the earth. We have filtered out a part of imaging area of the CCD chips where calibration isotopes irradiate. These area corresponds to two corners of each chip. Hot pixels of the chips are also removed by following the standard procedure. These criteria are summarized in Table 6.1. We have not filtered out the data with the Cut Off Regirity (COR).

6.1.2 Imaging analysis

In order to have a perspective view on the X-ray emission observed with Suzaku in the Galactic Center Region, we have created the mosaic X-ray images of the Galactic center region with the following steps. First, ”raw images” are obtained from the individual CCD censors, XIS-0,1,2,3, with an energy range of 2–10 keV. The image is smoothed

55 Table 6.1: Criteria of data analysis for the XIS.

GRADE 0:0 2:4 6:6 COR no selection ELV NTE > 5◦ ELV DYE > 20◦ Time after SAA > 436 sec

1.4 HXD-PIN 1.2 +0.5° FOV 1.0 Sgr D 0.8 Sgr B 1E 1743 Sgr C 0° Sgr A* 1E 1740 0.6 KS 1741 0.4

-0.5° 0.2 +1.5° +1.0° +0.5° 0° 1A 1742 -0.5° -1.0° -1.5° 0

East Region Center Region West Region

Figure 6.1: XIS image (2−10 keV). with a Gaussian of σ = 1.2′. The Non-Xray Background (NXB) is not subtracted in the imaging analysis. In order to compensate the vignetting effect, these images are then divided by ”flat images” calculated by a program called xissim (Ishisaki et al. 2007). In the program, an image expected from diffuse emission with no spacial variation (flat image) is simulated by a Monte Carlo method based on the detector response of XRT and CCDs. Finally, the images are corrected with effective exposures of the individual observations and mapped on the sky coordinate as shown in Figure 6.1. In Figure 6.1, extremely bright X-ray emission is seen in a region between ”Sgr A*” and ”1E 1743.1-2843”. The bright X-ray emission is so-called the Galactic diffuse emission ( Chapter 2 ) As clearly seen in the figure, the X-ray emission from Galactic center region consists of the diffuse emission and bright X-ray point sources. Diffuse emission are classified into two, one is the emission centered at the Galactic Center and the other is the emission from other diffuse structures probably associate with giant molecular clouds or SNRs. Since the FOV of HXD-PIN is larger than that of XIS as shown in Figure 6.1, we

56 need to study special variation of spectra in the FOV of HXD-PIN, before comparing two spectra. For convenience, we have divided all the observations into four regions, which roughly corresponds to the FOV of HXD-PIN. The each region is shown in Figure 6.1. The first region is the observations of the center region whose the galactic longitude is l < 0.3◦ and the galactic latitude b < 0.2◦. In this region, a non-thermal emission has been suggested by Koyama et al. (2007c) with the XIS spectra. The second region is the west region whose the galactic longitude is l < 0◦ except for the center region, which has no bright structures except for several point sources. The third region is the east region whose the galactic longitude is l > 0.4◦, where the giant molecular clouds ”Sgr B2” exists. Since Sgr B2 is known X-ray emitter, both the Galactic diffuse emission are probably contaminated by the emission from Sgr B2. The forth region is the region whose detected flux is filled with emission from bright point sources. We analyzed the XIS images for the each region. Definitions of the above four regions are listed in Table 6.2

57 Table 6.2. Deﬁnition of the 4 regions.

Region Sequence Positiona Start Expb number l(◦) b(◦) (UTC) [ks]

Center 100027020 -0.247 -0.046 2005-09-24 14:16:00 33.4 100037010 -0.247 -0.046 2005-09-29 04:25:00 36.7 100027010 0.057 -0.074 2005-09-23 07:07:00 34.9 100037040 0.057 -0.074 2005-09-30 07:41:00 33.0 100048010 0.057 -0.074 2006-09-08 02:11:00 56.1 West 501053010 -1.833 -0.000 2006-10-10 21:18:00 18.5 501052010 -1.500 0.000 2006-10-10 06:44:00 14.9 501051010 -1.167 -0.000 2006-10-09 13:39:00 19.4 500019010 -1.091 -0.041 2006-02-23 10:50:00 11.0 501057010 -1.167 -0.333 2006-10-11 10:06:00 17.5 501049010 -1.167 0.333 2006-10-08 10:19:00 16.2 501048010 -0.833 0.333 2007-03-11 19:04:00 22.3 501047010 -0.500 0.333 2007-03-11 03:55:00 17.8 501046010 -0.167 0.333 2007-03-10 14:43:00 22.9 East 501040010 0.607 0.072 2006-09-21 17:21:00 49.8 100037060 0.637 -0.095 2005-10-10 12:07:00 65.9 501039010 0.780 -0.160 2007-03-03 12:05:00 84.6 100037070 1.000 -0.100 2005-10-12 07:05:00 8.6 501059010 1.167 0.000 2007-03-15 18:55:00 49.6 501058010 1.300 0.200 2007-03-14 05:00:00 46.8 501060010 1.500 0.000 2007-03-17 05:06:00 50.2 Bright 501050010 -0.833 0.000 2006-10-09 02:19:00 17.5 Sources 501056010 -0.833 -0.333 2007-03-13 15:40:00 23.1 500018010 -0.569 -0.093 2006-02-20 12:30:00 43.8 501055010 -0.5 -0.333 2007-03-12 23:58:00 19.9 100027040 -0.446 -0.067 2005-09-24 12:40:00 1.6 100037030 -0.446 -0.067 2005-09-30 06:05:00 2.5 100027030 -0.441 -0.389 2005-09-24 11:05:00 1.7 100037020 -0.441 -0.389 2005-09-30 04:29:00 2.1 501054010 -0.167 -0.333 2007-03-12 08:09:00 21.5 501008010 -0.154 -0.191 2006-09-26 14:14:00 102.6 501009010 -0.074 0.178 2006-09-29 21:25:00 43.6

58 Table 6.2—Continued

Region Sequence Positiona Start Expb number l(◦) b(◦) (UTC) [ks]

100027050 0.328 0.010 2005-09-25 17:27:00 1.7 100037050 0.328 0.010 2005-10-01 06:21:00 2.0 500005010 0.428 -0.117 2006-03-27 22:40:00 60.4

aFOV (XIS-nominal) center position.

bEﬀective exposure of the HXD-PIN data.

6.1.3 Analysis of the XIS spectra We have extracted the XIS spectra for each region. Energy responses of the XIS spectra are created with xisrmfgen in HEASOFT 6.3.1 (http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/), which gives an energy response of XISduring an observational period. Ancillary Response File (ARF) are created with xissimarfgen also in HEASOFT 6.3.1. The ARF corrects the eﬀect of vignetting of XRT. In the spectral analysis, the NXB is estimated by the observational data base of the night earth (Koyama et al. 2007a). In the spectral ﬁtting, we have used xspec with a version of 11.3.2 (Arnaud 1996).

Center Region

Koyama et al. (2007c) have analyzed the XIS spectra of the center region and suggested that the non-termal emission exist in the Galactic diffuse emission (section 2.4) . Figure 6.2 show the XIS spectrum summed up the observational data of ”100027010”, ”100027020”, ”100037010”, ”100037040” and ”100048010”. The spectrum of the left panel is fitted with ”absorption × ( a thermal plasma model + 3 lines ). The absorption is the interstellar one and assumed to be 6 × 1022 /cm2 after Koyama et al. (2007c). We have used the thermal plasma model as ”apec” in xspec (http://hea- www.harvard.edu/APEC/). The model is an emission spectrum from collisionally-ionized diffuse gas based on the built-in atomic data. The spectrum consists of line emissions from the ionized elements and a continuum component from a thermal bremsstrahlung by electrons. The added 3 lines are fluorescence X-ray emissions from neutral Fe and Ni and are considered to be from the molecular clouds in the Galactic center region. It is suggested that the X-ray emission from the past activity in the Galactic center illuminate the neutral Fe and Ni in the molecular clouds and electrons in these atoms are exited (Murakami et al. 2001; Koyama et al. 2007b) . The energy of the individual lines are

59 6.4 keV (Fe I K-α), 7.0 keV (Fe I K-β) and 7.5 keV (Ni I K-α). The obtained spectrum can not be explained by the above model and another component of a power-law model is necessary as mentioned in section 2.4. The right panel of Figure 6.2 shows the result of the ﬁtting with the above model added a power-law model. The power-law model is shown as dN ∝ ϵ−Γ (6.1) dϵ where Γ is a photon index. The spectrum can be explained by the combination of these models. Hence, we deﬁne the model ”absorption × ( a thermal plasma model + 3 lines + power-law)” as ”Model A” in this thesis. Our interest for the XIS spectra is the continuum component extending to the hard X-ray band. The parameters which determine the continuum component are the temper-

ature of the hot plasma (kT), the Fe abundance (ZFe) which determines the equivalent width of Fe lines to the thermal bremsstrahlung and the photon index of the power-law component (Γ) . These best ﬁt parameters are listed in Table 6.3. The temperature is precisely determined to 6.5 ± 0.1 keV, while the Fe abundance and the photon index are not determined with only the XIS spectrum as shown by Koyama et al. (2007c). Hence, the HXD-PIN spectra is necessary in order to investigate the power-law component.

Table 6.3: The best ﬁt parameters of the XIS spectra summed with the center, west and east region.

a 2 Region kT [keV] ZF e Γ Flux5.5−11keV χν(ν) +0.09 +2.0 +0.5 Center 6.47−0.05 1.2−1.2 1.40−0.5 2.9 1.19 (933) +0.22 +8.0 +3.9 West 6.30−0.14 1.1−1.1 1.37−4.4 0.4 1.24 (290) +0.10 +1.0 +2.6 East 6.53−0.11 1.2−0.2 1.44−2.5 0.7 1.28 (737)

aThe unit is 10−10 erg/cm2/sec

West and East Region We verify if the spectra of the west and east region are explained by Model A. Before the spectral analysis, we have analyzed the image of the west and east region and removed point sources (Figure 6.3, 6.4). Some parts are above 3 σ signiﬁcance level and removed with 3.5′ radius circle. The individual spectra are shown in Figure 6.5.The spectra re- semble each other and that of the center region. However, the statics of the individual spectra are not good, and therefore, we have analyzed the spectra summed for the individual regions. Figure 6.6 shows the summed spectrum of the west region. We have

60 a ihteefie aaeesadtecniumcmoet r h aeeeyhr in fitted everywhere be same can the spectra are the region. components all center continuum that Galactic the confirmed the and is parameters It fixed = 6.4. these (kT Table (2007c) with in al. listed the et are Koyama assume parameters by we photon fixed shown Hence, the region center spectra. or Z the XIS keV, abundances of Fe regions 6.5 those three the be the either to of determine parameters fitting three to spectral able the not with are indices We true shape. is same same the The the region. region. of east center parameters explained the the the of be that of spectrum confirmed can those also the spectrum with is in it the consistent And, are that region. component confirmed center continuum and the like A A Model Model with with spectrum the fitted h nti 10 is unit The ehv eie ftecniumcmoet ntesetao h he ein are regions three the of spectra the in components continuum the if verified have We χ normalized counts/sec/keV aaeesaefixed. are parameters Reduced 6.4: Table (right). model power-law plasma a thermal a added with plasma fitted thermal region a center and the (left) of model spectra XIS The 6.2: Figure −5 0 150−3 0.01 0.1 1 eink kV Z [keV] kT Region etr65(xd . fie)14 (fixed) 1.40 (fixed) 1.2 (fixed) 6.5 Center et65(xd . fie)14 (fixed) 1.40 (fixed) 1.2 (fixed) 6.5 West at65(xd . fie)14 (fixed) 1.40 (fixed) 1.2 (fixed) 6.5 East Fe − 1 = 10 erg/cm . ,Γ=1 = Γ 2, c h 2 a /sec n n e l

e n χ . e r ) h reduced The 4). g 2 y

( k fte3rgo pcr te ihMdlAwoemain whose A Model with ﬁtted spectra region 3 the of e V ) e F 1 0 61 χ

2 χ normalized counts/sec/keV

ftetreXSsetaﬁtdwt these with ﬁtted spectra XIS three the of −3 Γ −2 0 2 10 0.01 0.1 1 Flux 5 0.7 0.4 2.9 . c 5 h − a n n 11keV e l

e n e r g y

( a k e V ) .7(740) 1.27 .4(293) 1.24 (936) 1.24 χ ν 2 ( ν ) 1 0 Figure 6.3: XIS images of the 9 observations in the west region. Contours are 2, 3 and 4 σ signiﬁcance levels.

62 Figure 6.4: XIS images of the 7 observations in the east region. Contours are 2, 3 and 4 σ signiﬁcance levels.

63 1 1 . . 0 0 V V e e k k / / c c e e s s 1 1 / / 0 0 s s . . t t 0 0 n n u u o o c c d d e e z z i i l l a a m m r r o o 3 3 − − n n 0 0 1 1 4 4 − − 0 0

1 10 1 10 channel energy (keV) channel energy (keV)

Figure 6.5: The XIS spectra of the west (left) and east (right) region . 1 . 0 V e k / 1 c 0 e . s 0 / s t n u o c

d 3 e − z 0 i l 1 a m r o n 4 2 χ 0 2 − 4 − 10 channel energy (keV)

Figure 6.6: The XIS spectrum summed with the 8 spectra of the west region.

64 Figure 6.7: The XIS spectra ﬁtted with the model of the west region.

65 Figure 6.8: Same as Figure 6.7

66 bandfnto sdana ledse iei iue69 h cl egto 0.15 of height scale The 6.9. of Figure points in data line 3 the dashes by blue obtained a is as drawn is function obtained h eie itiuinmdli ieetwt h . e eln ue distribution fluxes line Fe keV 6.7 the with different is model distribution derived The obtained is distribution 0.8, brightness = surface the C X-ray in 9.3, soft Hence, = The FOV. (1998). G XIS as the Maeda by in shown points degree data calculate 0.5 ” we of as length fitting, shown scale spatial is the functions for The negligible not (1998). Maeda after exp( model fitting a as figure, the this center. in In the shown toward As clearly excluded. 6.9. increases are distribution Figure sources the point shown bright as by contaminated longitude of significantly points Galactic plotting data distribution the the brightness along figure, surface obtain factors X-ray we Second, soft normalization the 6.8. the ”normalization derive Figure as we and spectra Finally, 6.7 individual region Figure parameters)”. west the in the for of shown factors spectra are normalization XIS Z model the The keV, above freed. 6.5 the all with = are with summed) lines (kT fitted neutral fixed (not three are observation the parameters of individual calculate parameters main the The three to of the have spectra whose we brightness A the HXD-PIN, surface Model fit X-ray and soft we XIS the First, obtain of we distribution. spectra Therefore, FOVs. band two wide the between the differences X-ray analyze to Soft order of In Distribution Brightness Surface 6.1.4 nodrt oe h itiuin eaota xoeta lscntn function constant plus exponential an adopt we distribution, the model to order In | b o h I ue.Tebakdt onsso h . e iedistribution line keV 6.7 the show points 2007c). data al. black et The (Koyama observations fluxes. the XIS of the data for The (red). fluxes XIS of the of Distribution 6.9: Figure | /b < b 0 + ) 0 C . 1 .Temdli o-ieradtesz fXSFVo 0.3 of FOV XIS of size the and non-linear is model The ”. ) ◦ r lte.Tebu iesostersl ftesailfitting spatial the of result the shows line blue The plotted. are Relative soft X-ray intensity l 0 0 = | b | χ < . 33 2 0 rmdffrne fitgasbtentemdladthe and model the between integrals of differences from ◦ . 1 and ◦ l r lte,adtedt onso h observations the of points data the and plotted, are = b − 0 1 0 = . 2 67 ◦ . ntews eina hw nFur 6.10. Fiugre in shown as region west the in 15 ◦ n h tigerris error fitting the and , G Fe × × 1 = oe (fixed A Model exp( ∼ . ,Γ=1 = Γ 2, 0% The %. 20 ◦ −| × l | 0.3 /l 0 ◦ . ) 4). × is ◦ Figure 6.10: Distribution of the XIS fluxes of the west region. Black circle; b = 0.3◦. Red square; b = 0◦. Black cross; b = −0.3◦

(Koyama et al. 2007c) drawn as black data points in Figure 6.9 by ∼ 0.1◦. However, the difference is negligible for the size of the HXD-PIN FOV. In order to analyze the wide band spectra of XIS and HXD-PIN for extended emission (section 6.4), it is necessary to consider the difference of the FOVs between two detectors. For the XIS spectra, we make the ARF with xissimarfgen (section 6.1.3) assuming the emission is uniform in the FOV and the size of the emission is 20′ radius circle (1257 min2). While, for the HXD-PIN spectra, we utilize the response file created assuming the uniform emission whose size is 2◦ × 2◦ (14400 min2). In addition, the cross normalization between the XIS and HXD-PIN spectra for point sources is calibrated to be 1.13 (Kokubun et al. 2007). Hence, in the actual analysis, we may adjust the difference of the FOVs by the ratio of 14400/1257 = 11.5 multiplied by 1.13, that is 12.9, for the perfectly uniform emission. However, another adjustment (hereafter ”constant factor”) is necessary since the Galactic diffuse emission is not perfectly uniform in the HXD-PIN FOV. In the wide band spectral analysis, we utilize products of 12.9 by the constant factor as the adjustment. We calculate the constant factors as bellow. First, we assume the FOV of XIS and HXD-PIN as a rectangular solid with a base plane of 0.3◦ × 0.3◦ (18′ × 18′) and the HXD-PIN angular response of the triangle (described in chapter 4) respectively. Next, we calculate overlapping volumes of the soft X-ray surface brightness distribution and the FOVs of XIS and HXD-PIN respectively for the each position. In this calculation, we consider the roll angle of the satellite as 40◦ for the galactic longitude which is typical angle of the actual observational attitude. The calculated distribution of XIS is drawn as a blue solid line in Figure 6.9. And, The calculated distribution of HXD-PIN is drawn in the middle panel of Figure 6.11. The constant factor is calculated from the ratio of the rounded XIS distribution to the rounded HXD-PIN distribution, and shown in the

68 Figure 6.11: Top panel; distribution of the rounded XIS distribution. Middle panel; distribution of the rounded HXD-PIN distribution. Bottom panel; distribution of the constant factor which means diﬀerence of the FOVs between XIS and HXD-PIN. The blue solid line and the black dashes line are calculated along b = 0◦ and b = 0.3◦ respectively. bottom panel of Figure 6.11. We also calculate the constant factor distribution along b = 0.3◦. These constant factors are utilized in section 6.4.

69 6.2 HXD-PIN Data Analysis

6.2.1 Data Reduction We experienced unstable operation of 16 PIN diodes installed in Well-counter units, W00–W03, hereafter W0. The bias voltage for these diodes was reduced to 400 V from the nominal voltage of 500 V on 26 May 2006. On 2006 October 4, 16 more PIN diodes (the PIN diodes in the Well-counter units of W10–W13, that is W1) was reduced to 400 V for the same reason. The reduction of bias voltage caused reduces of the thickness of depletion layers in PIN-Si, hence leads to decreases of the effective area and their energy responses. Furthermore, we need to accumulate the Non X-ray Background (NXB) spectrum under the new condition, in order to subtract backgrounds properly. Hence, we prepared response matrix and NXB model for 400V-bias. For HXD-PIN, we also used Suzaku data sets processed by the Suzaku data processing pipeline version 2.0. together with the calibration data set of version 2007-09-15. We filtered out the data obtained during passages through the South Atlantic Anomaly. And the data was filtered with the elevation angle to the Earth’s limb below 5◦. Since the repeatability of the NXB is not assured in the orbits of COR < 6 GeV, we also filtered the data in the orbits.

6.2.2 NXB Modeling The NXB models of HXD-PIN are constructed using the database of earth occultation data, whose average exposure was ∼ 15 ks per day during half a year since the launch (Watanabe et al. 2007; Kokubun et al. 2007). In the modeling, the counter of upper discriminator of the PIN diodes (PIN-UD) is used as a monitor of real-time particle ﬂux. The threshold level of the PIN-UD corresponds to ∼ 90 keV, and its counting rate can be regarded as the number of cosmic-ray charged particle penetrating the device. Since the PIN diodes are embedded in the thick BGO shields, only protons above ∼ 100 MeV can trigger the PIN-UD. Figure 6.12 shows typical lightcurves of the total PIN-UD count rate, the event rate of PIN and GSO, and the time variation of COR during a blank sky observation. The sharp peaks of ∼ 10000 counts s−1 in the PIN-UD lightcurve indicate SAA passages. Although the real time PIN-UD count rate can be used as a good indicator of the NXB component which directly correlate with COR, another NXB component which is correlated with T SAA cannot be modeled by simply using PIN-UD. As mentioned above, this component is relatively minor compared to that correlated with COR. However, it is by no means negligible as shown in the lightcurve. In the NXB model, this component is taken into account using a parameter deﬁned as, ∫ ( ) t t0 − t PIN-UDbuildup(t) = PIN-UD(t0) exp dt0. (6.2) −∞ τ

70 Figure 6.12: A typical lightcurve of the PIN-UD summed over the 16 units together with that of events from PIN and GSO, and the COR, obtained from ∼ 1.5 days observation of a blank sky ﬁeld (Kokubun et al. 2007).

Various values of time constant τ was tried between 5000–10000 s and τ = 8000 s was selected because this value gives the best correlation between the model and the earth

occultation data. A typical lightcurve of PIN-UD and PIN-UDbuildup are shown in Fig- ure 6.13. The NXB database was constructed by accumulating and sorting the earth occultation data by these two parameters. The background for each observation is estimated by picking up the spectra from the database sorted by the two parameters. The reproducibility of the NXB model is examined by Mizuno et al. (2006). They show that the current accuracy of the NXB model is about 5%.

6.2.3 Spectral Analysis

We analyzed the HXD-PIN spectra from all the observation. The exposures of the spectra were corrected with counts of ”pseudo trigger events”. Since the pseudo trigger is randomly vetoed by normal events, we can accurately estimate the dead time fraction, by comparing the total number of injected pseudo pulses and those actually recorded in the data. The HXD-PIN spectrum from each pointing was constructed and compared with the background model estimated for the each observation period. Since the NXB model does not include the contributions from CXB, simulated spectrum of CXB was added to the NXB model. Here, the CXB spectrum was modeled as ( ) ( ) dN ϵ −1.29 ϵ = 7.9 × exp − photons s−1 keV−1 cm−2 str−1, (6.3) dϵ 1 keV 41.13 keV

which is based on the HEAO-1 observations in 1970’s (Gruber et al. 1999). However, it is known that large absorption of ∼ 1022 − 1023 nH cm−2 exists in the Galactic center region. Therefore we estimated the CXB obtaining the absorption of 6 × 1022 nH cm−2 referring Koyama et al. 2006, although an eﬀect of the absorption is less than 1% and

71 Figure 6.13: Lightcurves of PIN-UD (top) and PIN-UDbuildup (bottom) (Tanaka 2007).

negligible above 10 keV. From this model, the contribution from the CXB flux to the detected counting rate can be estimated to be ∼ 5% of the NXB. Figure 6.14 and Figure 6.15 shows an explanation of the HXD-PIN spectra and lightcurves respectively. And, Figure 6.16 - 6.24 show spectra and lightcurves extracted from all the observation in this thesis. These spectra are subtracted with the NXB model and the CXB. It is confirmed that HXD-PIN detects bright hard X-ray of 1–4 times of the NXB. Figure 6.25 shows a distribution of the NXB-subtracted counting rate in 16– 40 keV. In the figure, large signals were detected in all the observations and concentrated on the center. Since a shape of the distribution is broader than the angular response of HXD-PIN, it is not able to be explained with contaminations of a few hard X-ray sources at least. Figure 6.26 shows a distribution of the each PIN’s counting rate also subtracted the NXB in 16–40 keV. The structure of the distributions are divided to two shapes roughly; sharp and gradual these. The plots of the observations which shape sharp curves were possibly influenced by bright point sources. The variability between different color plots possibly shows the fluctuation of a point source. While, gradual curves possibly mean the diffuse emission or superposition of many hard X-ray sources.

6.2.4 Spectral Fitting with a Power-law Model The spectra extracted from the observations whose exposures are above 5 ksec are fitted with a power-law model. The absorption is assumed to be 6 × 1022 nH cm−2, again. We perform the spectral fitting with an energy range of 14–40 keV. In the spectral fitting, we use xspec with a version of 11.3.2. The best-fit parameters are summarized in Table 6.6.

72 Raw spectrum Error = Statical error of the detected events

NXB Error = Statical error of the Montecarlo simulation (Negligible because of enough statics)

Detected spectrum (Celestial signals) Error = The same as the error of the raw spectrum Reproducibility of NXB

Uncertainity of the NXB is estimated with the reproducibility (Almost negligible because of the intense hard X-ray ﬂuxes in this thesis)

Figure 6.14: Explanation of the HXD-PIN spectrum.

NXB lightcurve catchs up the ﬂuctuation by SAA SAA SAA

Raw ﬂuxes

NXB

Flattened Detected ﬂuxes

Figure 6.15: Explanation of the HXD-PIN lightcurve.

Table 6.5: Criteria of data analysis for the HXD-PIN.

COR > 6 GeV ELV > 5◦ Time after SAA 436 sec

73 Table 6.6: Summary of power-law model ﬁtting to the HXD-PIN spectra.a

2 Sequence Γ Count Rate χν (ν) number (14–40 keV) [mCrab]a +0.32 ± 501053010 2.37−0.30 1.52 0.11 0.78 (26) 501052010 1.77 ± 0.14 4.20± 0.14 0.94 (27) 501051010 1.74 ± 0.04 18.47± 0.18 1.37 (45) +0.28 ± 500019010 2.81−0.27 2.70 0.15 1.54 (19) 501057010 1.94 ± 0.04 18.01± 0.19 0.75 (43) +0.15 ± 501049010 2.07−0.14 4.13 0.13 1.15 (29) 501048010 1.68 ± 0.05 11.53± 0.14 0.81 (44) 501047010 1.89 ± 0.07 9.37± 0.15 1.39 (37) 501046010 2.36 ± 0.05 14.90± 0.15 0.64 (44) 501050010 1.77 ± 0.03 27.05± 0.21 0.87 (48) +0.03 ± 501056010 1.77−0.02 32.33 0.20 1.15 (54) 500018010 2.16 ± 0.03 17.58± 0.11 1.10 (55) 501055010 2.20 ± 0.03 34.59± 0.22 0.95 (51) 100027020 2.66 ± 0.03 18.09± 0.13 1.17 (52) 100037010 2.78 ± 0.03 16.97± 0.12 1.11 (52) +0.04 ± 100027010 2.85−0.03 17.15 0.13 0.83 (50) 100037040 2.82 ± 0.04 17.06± 0.13 0.69 (50) 100048010 2.70 ± 0.03 15.27± 0.10 1.35 (57) 501054010 2.49 ± 0.03 25.14± 0.19 1.09 (47) 501008010 2.52 ± 0.02 19.32± 0.08 1.63 (67) 501009010 2.61 ± 0.05 9.67± 0.10 1.08 (51) 500005010 2.26 ± 0.02 34.63± 0.13 1.41 (65) 100037060 2.42 ± 0.04 8.49± 0.07 1.22 (56) 501040010 2.65 ± 0.06 7.01± 0.08 1.04 (51) 501039010 2.44 ± 0.05 6.25± 0.06 0.85 (58) 100037070 2.37 ± 0.20 4.18± 0.18 1.77 (16) 501059010 2.32 ± 0.10 3.55± 0.07 1.03 (49) +0.14 ± 501058010 2.27−0.13 2.50 0.07 0.88 (46) 501060010 2.26 ± 0.15 2.15± 0.07 0.86 (47)

aErrors represent 90% conﬁdence. b1 mCrab corresponds to 0.045 counts/sec in the energy range of 14–40 keV.

74 Figure 6.16: The HXD-PIN spectrum (left) and lightcurve (right) in the range of 14–40 keV.

75 Figure 6.17: Same as Figure 6.16

76 Figure 6.18: Same as Figure 6.16

77 Figure 6.19: Same as Figure 6.16

78 Figure 6.20: Same as Figure 6.16

79 Figure 6.21: Same as Figure 6.16

80 Figure 6.22: Same as Figure 6.16

81 Figure 6.23: Same as Figure 6.16

82 Figure 6.24: Same as Figure 6.16

83 Figure 6.25: Distribution of hard X-ray ﬂux extracted from observational data in the energy range of 16–40 keV. Red plots are data of the observations with b < 0.2◦ and black plots are that of b > 0.2◦.

Figure 6.26: Distribution of hard X-ray ﬂuxes of each PIN extracted from observational data in the energy range of 16–40 keV.Diﬀerence of plotting color means terms performed the observations. Black; 2005/09–2005/10. Red; 2006/02–2006/03. Green; 2006/09–2006/10. Blue; 2007/03.

84 PIN’s FOV (full width)

AX J1749.2-2725 GRS 1741.9-2853 1E 1743.2843 AX J1749.1-2733 KS 1741-293 IGR J17475-2822 1E 1740.7-2942 1A 1742-288 1A 1742-294 IGR J17497-2821

SLX 1744-299/300

Figure 6.27: Hard X-ray sources detected by the IBIS with signiﬁcances of above 3σ.

6.3 Estimation of Contaminations from Known Hard X-ray Sources

6.3.1 Known Hard X-ray Sources in the Galactic Center Region

The INTEGRAL onboard Imager (IBIS) and the Burst Alert Telescope (BAT) onboard Swift cover almost whole the energy range of the HXD-PIN. INTEGRAL is monitoring the Galactic center region about once every 1–4 days under the Galactic Bulge Monitoring Program (Kuulkers et al. 2007). While, BAT has an extreme broad FOV of one-sixth as large as all sky, and therefore, is able to monitor the Galactic center region almost at least once every a few days even if the center of the FOV is not directed at the region. In addition, both the detectors covered terms when Suzaku observed the Galactic center region. However, in the energy range of 20–60 keV, a sensitivity of the IBIS onboard INTEGRAL is better by 2–3 times than the BAT with the same exposure. We, therefore, have utilized the IBIS as the monitor of hard X-ray sources in this thesis. The hard X-ray sources in the Galactic center region detected with above 3 σ signif- icances by IBIS are shown in Figure 6.27 with the FOV (full width) of the HXD-PIN. It is clear that it is impossible to avoid contaminations from these sources to the FOV even with the narrowest FOV among those of collimator-type hard X-ray detectors. The names, positions, objective types, typical ﬂuxes and typical time scale of ﬂuctuations are listed in Table 6.7. Fluxes of some sources are several 10 mCrab and it is not negligible compared with those detected by the HXD-PIN as described in section 6.2. In almost all

85 the observation, there is at least one source which contaminates in the HXD-PIN FOV. Therefore it is crucial to estimate contributions of the sources to the HXD-PIN ﬂuxes in order to study the non-thermal emission in the Galactic Center region. We have took account of these known 12 sources in this thesis and estimated the contributions to the HXD-PIN ﬂuxes.

Table 6.7: Summary of the hard x-ray sources in the Galactic Center.

Source name Position Averaged ﬂux Time scale Type l(◦) b(◦) [mCrab] [days] AX J1749.2-2725 1.70 0.12 ∼3 — HMXBa AX J1749.1-2733 1.59 0.05 ∼4 — HMXBa IGR J17497-2821 0.95 -0.45 ∼13 ∼10 BHa IGR J17475-2822 0.61 –0.06 ∼5 — Molecular Cloud? 1E 1743.1-2843 0.26 –0.03 ∼3 — LMXBa 1A 1743-288 0.21 –0.24 ∼5 — LMXBa IGR J17456-2901 –0.06 –0.05 ∼4 — BHa(SgrA*) GRS 1741.9-2853 –0.05 0.12 ∼ 5 — LMXBa KS 1741-293 –0.43 –0.09 ∼6 ∼30 LMXBa 1A 1742-294 –0.44 –0.39 ∼17 ∼50 LMXBa SLX 1744-299/300 –0.74 –0.91 ∼8 — LMXBa 1E 1740.7-2942 –0.88 –0.11 ∼40 ∼100 BHa

aHMXB; High Mass X-ray Binary. BH: Black Hole Candidate. LMXB; Low Mass X-ray Binary.

6.3.2 Flux Estimation of the Hard X-ray Sources with IBIS We have estimated the fluxes of the hard X-ray point sources basically with the IBIS lightcurves. And, we define the estimation with the IBIS lightcurves as ”the IBIS method”. Lightcurves of the hard X-ray sources detected by IBIS are shown in Fig- ure 6.28 and Figure 6.29. We have utilized these lightcurves to estimate the fluxes of the sources in the FOV of the HXD-PIN. We show the term, when the Suzaku observation took place, as blue lines in the figure. We can see clear time variation during the observation with Suzaku. Figure 6.30 shows an example of a blow-up of lightcurves of ”1A 1742-294”. Three sets of monitoring were performed during 32 days when Suzaku observed, which is the monitor every 4 days roughly. In the cases of almost all the other terms and sources, the same is true or they are more frequent. Errors of the lightcurves are constant roughly and ∼ 1 counts/sec which is equivalent to 4 mCrab.

86 Figure 6.28: Lightcurves of the hard X-ray sources with IBIS (20 - 60 keV).

87 Figure 6.29: The same as Figure 6.28

88 Figure 6.30: Enlarged lightcurve of 1A 1742-294.

Figure 6.31: Lightcurves of 1A 1742-294 ﬁtted with a linear function. MJD of the blue circle means estimated ﬂux during the observation of ”501008010”

89 We estimated the fluxesby fitting the lightcurves with a linear function. An example of the fitting is shown in Figure 6.31. In the figure, Modified Joulian Day (MJD) of the blue dots is the day of Suzaku observation whose sequence number is ”501008010”. In fitting the lightcurves, we have selected the fitting plots to take in at least 3 plots both before and after the target MJD. We have set the fitting range ±1 day at first, and have expanded the range by 1 day if the number of the plots are fewer than 3. The fitting errors are ∼ 0.3 counts/sec which is equivalent to 1 mCrab. In order to verify if a linear function is adequate to the lightcurve fitting, we have calculated structure functions (SFs) from the individual lightcurves. Typical time scaleof time variability of individual sources can be estimated from the break point of the SFs. If the time scales are smaller than frequencies of the monitoring which are several days, the fitting with the linear functions are not adequate. We use the formalism described in Simonetti, Cordes, & Heeschen (1985). One modifications is that since photon statistics dominate the error in X-ray data, we use a continuous weighting factor proportional to its significance of each data point instead of taking only either 0 or 1 as in Simonetti, Cordes, & Heeschen (1985). The definition of the 1st order SF for a light curve described as f(i) is then,

1 ∑ SF (τ) = w(i)w(i + τ)[f(i + τ) − f(i)]2, (6.4) N(τ) i where ∑ N(τ) = w(i)w(i + τ), (6.5) i

f(i) w(i) ∝ . (6.6) σf (i)

Here, w(i) is the weighting factor, and σf (i) is the 1 σ uncertainty of the data point f(i) in the lightcurve. The summations are made over all pairs. Figure 6.34 shows a schematic drawing of the ’typical’ SF for measured time series. In the figure, constant offset means white noise which is made from measurement noise and real fluctuations not depending on frequencies, which are not able to be resolved. The slope of τ β means red noise which is random walk, and the slope increases until typical time scale of the fluctuations. Figure 6.32 and Figure 6.33 show the structure functions extracted from the lightcurves of Figure 6.28 and Figure 6.29. We have confirmed that, in all the case of the sources but ”IGR J17497-2821”, the typical time scales are above 10 days, or the SFs’ have no typical time scale and then are shown as white noises in whole frequencies. Therefore, the monitoring frequencies of IBIS are enough to follow up the fluctuations of the sources and the fitting with a linear function is adequate for these light curves. In the case of ”IGR

90 Figure 6.32: Structure functions extracted from the light curves of Figure 6.28.

91 Figure 6.33: Structure functions extracted from the light curves of Figure 6.28.

92 log SF(τ)

White noise

Typical time scale log τ

Figure 6.34: Schematic drawing of the ’typical’ structure function for measured time series.

J17497-2821”, the SF have mostly formed by one flare at the MJD of ∼54000 as shown in Figure 6.29. This effects only the observation of ”500005010”. The IBIS lightcurve of ”IGR J17497-2821” can not follow up the real flux of the source in the observational term of ”500005010”. We have evaluated other uncertainties of the flux estimations by following methods. First we have excluded a real data point in the lightcurves, and we have performed the fitting the same as described above using the data points around the excluded data point. Second, we have compared the estimated count rate from the fitting with that of the real data point. In the fitting, however, we have also excluded the data points within 1 day from the MJD of the target data point. Differences between the count rate from the fitting and of the real data point are distributed with root mean squares of ∼ 1 counts/sec, which are comparable to statical errors of the actual√ data points. These are also comparable to the white noises in the SFs’ which is ∼ SF (τ)/2. These uncertainties mean unknown fluctuations within the interval of the IBIS lightcurves. We have adopted the uncertainty as a systematic error and named as the short fluctuation uncertainties. That is, finally, the errors are shown as error = (fitting error)±(short fluctuation uncertainty). The short fluctuation uncertainty is dominant in this error, and therefore the error is estimated to be ∼ 1 count/sec, 4 mCrab. In order to subtract the contaminations of the hard X-ray sources from the HXD- PIN fluxes using the IBIS fluxes, it is necessary to obtain a factor between counting rates of HXD-PIN and IBIS. We have obtained the factor using the HXD-PIN spectrum and the IBIS flux of the Crab. The analyzing process of the HXD-PIN spectrum is the same as that of the Galactic center observations as described in section 6.2. We have performed a spectral fitting for the Crab spectra adopting a power-law model in the

93 energy range of IBIS (20–60 keV), and obtained a photon index of 2.12 ± 0.01 and a flux of 1.39(±0.01) × 10−8 erg/cm2/sec at 90% confidence. While, the IBIS counting rate of the Crab is 240 counts/sec in 20–60 keV (Kuulkers et al. 2007). Therefore, we have defined the flux of IBIS of 240 counts/sec as 1.39×10−8 erg/cm2/sec when the photon index of the source is 2.12 in this thesis. A utilized energy response of IBIS is released by the INTEGRAL Science Data Centre, whose version is ”6.5”. In order to obtain the HXD-PIN spectra subtracted the contaminations of the sources, we need to assume shapes of the source spectra because we are not able to know the spectra with only the IBIS. Since dependencies of effective area on energy of between HXD-PIN and IBIS are different with each other, it is necessary even to obtain just the HXD-PIN fluxes subtracted the contaminations in IBIS’s energy range of 20–60 keV. Basically, we assume the spectra as simple power-law models from previous works for the individual sources as listed in Table 6.8.

Table 6.8: Summary of the assumed photon indices of power-law model for hard X-ray sources.

Source name photon index referecese AX J1749.2-2725 1.0 [1],[2],[3]a AX J1749.1-2733 1.0 [3],[4]a IGR J17497-2821 1.8 [5]a IGR J17475-2822 1.8 [6]a 1E 1743.1-2843 3.2 [7]a 1A 1743-288 2.4 [8]a IGR J17456-2901 3.0 [9]a GRS 1741.9-2853 1.6 [10]a KS 1741-293 2.3 [11]a 1A 1742-294 1.6 [12],[13]a SLX 1744-299/300 1.6 [12],[13]a 1E 1740.7-2942 1.5 [14]a

a[1]; Torii et al. (1998). [2]; Santangelo et al. (2006) . [3]; Sakano et al. (2002). [4]; Bird et al. (2002) [5]; ?. [6]; Revnivtsev et al. (2004). [7]; Del Santo et al. (2006). [8]; Natalucci et al. (2000). [9]; Belanger et al. (2006). [10]; Cocchi et al. (1999). [11]; De Cesare et al. (2007). [12]; Lutovniv et al. (2001). [13]; Tanaka et al. (1997). [14]; Del Santo et al. (2005)

94 6.3.3 Spectral Estimation of the Hard X-ray Sources with XIS

Some of the hard X-ray sources are also detected with XIS. The XIS spectra of the sources are more reliable than just adopting the spectral shape from the previous works. The fluxes in the energy band of HXD-PIN are calculated by extending the XIS spectra up to 40 keV. We define this flux estimation as the XIS method. Since the HXD-PIN’s FOV of full width is larger than that of the XIS by ∼15 times, there are few perfect simultaneous observations. Since the time fluctuations of the sources listed in the Table 6.7 are as long as a month except for one flare of ”IGR J17497-2821”, we are able to obtain the spectral shapes of the sources analyzing the data of the observations performed within abount a few weeks. Since we have performed 4 sets of the observations with terms of a month as described in chapter 5, we are able to obtain the spectral model of the target sources which are observed in the same terms. The sources detected with the XIS are listed in Table 6.9

Table 6.9: The list of the hard X-ray sources detected with XIS.

Source name Terma Sequence numberb ABCD AX J1749.1-2733 x x x o 501060010 1E 1743.1-2843 o x x x 100027050, 100037050 IGR J17456-2901 – – – – unavailable c KS 1741-293 o x x x 100027040, 100037030 1A 1742-294 o x x o 100027030, 100037020, 501055010 1E 1740.7-2942 x x o x 501050010

aA; 2005/09 – 2005/10. B; 2006/02 – 2006/03. C; 2006/09 – 2006/10. D; 2007/03. bSequence numbers of the observation in which the hard X-ray sources existed in the XIS FOV. cThe source is detected with the XIS, but the diﬀuse emission is too strong to extract the spectrum of just the source.

Six of twelve hard X-ray sources have been detected with the XIS. However, since ”IGR J17456-2901 (Sgr A*)” lies in very the center of the Galactic center region and the diﬀuse emission is as bright as this source, it is diﬃcult to extract the spectrum of just the source. Therefore, we have extracted the spectra of 5 sources for each observation terms. As described in chapter 5, since we have planed the strategies to select the 3 sources which are predicted to be especially bright in hard X-ray band (that is ”KS 1741-293”, ”1A 1742-294”, ”1E 1740.7-2942”) and monitor these sources with the XIS twice in the one term (which is ”A” term in the Table 6.9), we are able to obtain two observational

95 data every 3 sources. We have performed spectral fitting using the XIS spectra.Spectral parameters extracted from the fitting are summarized in Table 6.10. The spectra of the point sources are shown in Appendix. We have confirmed that the IBIS fluxes are consistent with the model determined by XIS spectra in the IBIS energy band (20 – 60 keV).

Table 6.10: Summary of the spectral ﬁtting for the bright point sources with the XIS spectra and the IBIS ﬂuxes.

a b c 2 Source name Term Model Column Density Parameters Flux χν [1022cm−2] kT[keV] Γ +1.2 +0.13 AX J1749.1-2733 D PL 6.7−1.2 1.18−0.13 0.04 1.23(43) +0.8 +0.1 1E 1743.1-2843 A1 BB 10.0−0.8 1.8−0.1 1.06 0.51(48) +0.9 +0.1 1E 1743.1-2843 A2 BB 11.6−0.9 1.8−0.1 1.56 1.10(88) +1.8 +0.20 KS 1741-293 A1 PL 18.6−1.7 2.15−0.19 0.53 0.62(20) +2.0 +0.2 KS 1741-293 A2 PL 16.2−1.9 2.0−0.2 0.25 0.51(14) +0.1 +0.4 1A 1742-294 A1 TB 5.7−0.1 8.6−0.4 6.33 1.05(409) +0.1 +0.4 1A 1742-294 A2 TB 5.5−0.1 9.5−0.4 5.58 1.06(565) +0.1 +0.8 1A 1742-294 D TB 6.1−0.1 18.3−0.7 2.96 1.00(1430) +0.6 +0.05 1E 1740.7-2942 C PL 13.8−0.6 1.52−0.05 1.61 1.08(488)

aA1; 2005/09/23 – 2005/09/25. A2; 2005/09/29 – 2005/10/01. C; 2006/09 – 2006/10. D; 2007/03. bPL;power-law. BB; black body. TB; thermal bremsstrahlung (bremss in XSPEC). cUnabsorbed ﬂux (2-10 keV) in units of 10−10 ergs/cm2/sec.

6.3.4 Flux Estimation of ”1E 1740.7-2942” We have estimated the flux of ”1E 1740.7-2942” on October 2006 with the multiple method, and compared the results with each other, which is important for the after analysis in section 6.4.The flux estimated with the IBIS method is 32 ± 4 mCrab (section 6.3.2). The flux with the XIS method is 31 ± 2 mCrab, and the error of the flux with the XIS method is estimated from the fitting errors of the photon index as shown in Figure 6.35 (Γ = 1.52 ± 0.05). And, we have tried to estimate the flux with the FC method and obtained 28 ± 3 mCrab. This flux is estimated from the observational data of ”501051010” in which the statical uncertainty is the smallest with the FC method. We have estimated the flux of ”1E 1740.7-2942” on October 2006 with another method. In this method, we utilize the three continuous observations along the Galac-

96 Figure 6.35: Flux estimation of ”1E 1740.7-2942” with the XIS method. tic plane at 0.3 degrees interval, ”501051010”, ”501052010”, ”501053010” or (l, b) = (−1.2, 0.0), (−1.5, 0.0), (−1.8, 0.0) (Figure 6.36). We are able to consider these as kinds of scanning observations and define as observation ”a”, ”b” and ”c” in this order. The contaminated known hard X-ray source in the observations of ”a” and ”b” is only ”1E 1740.7-2942” as mentioned after in section 6.4.2. And, in the HXD-PIN FOV of the observation ”c”, there are no contaminated sources, but only the hard X-ray diffuse emission, also as mentioned after. Summarized, the detected fluxes of ”a” and ”b” consists of the contaminations from the source and the hard X-ray diffuse emission, and that of ”c” only the diffuse emission. The three detected fluxes are 18.5 ± 0.2, 4.2 ± 0.1 and 1.5 ± 0.1 mCrab, respectively. And, the ratios of the angular response for the source in the observations of ”a” and ”b” are 0.47 and 0.07 respectively. When we define the flux of ”1E 1740.7-2942” and the diffuse emission as F and D, we obtain the following correlations; F × 0.47 + Da = 18.5 ± 0.2,F × 0.07 + Db = 4.2 ± 0.1,Dc = 1.5 ± 0.1. We assume the ratio of the diffuse emission from the XIS surface brightness distribution;

Da : Db : Dc = 2 : 1.3 : 1. Since an unknown value is only F , the flux of ”1E 1740.7- 2942” in above correlations, we are able to solve these equations and obtain the flux as 32 ± 2 mCrab. In this error, the statical uncertainties of the individual detected fluxes are dominant. The uncertainties are summarized in Table 6.11. We compare the fluxes estimated with the four method and confirmed that all the result are consistent with each other in the uncertainties as listed in Table 6.12.

97 Figure 6.36: Flux estimation of ”1E 1740.7-2942” with the three observations along the Galactic plane.

Table 6.11: Uncertainties of the ﬂux of ”1E 1740.7-2942” on October 2006 estimated with the three observations along the Galactic plane.

Comprising item Uncertainty Statical errors of the detected ﬂuxes 1.6 mCrab (dominant) Systematic error of angular response 1 mCrab Uncertainties of the assumed XIS ﬂuxes 0.7 mCrab Flux variation during the three observational term (2 days) 0.6 mCrab

Table 6.12: Comparison of the estimated ﬂuxes of ”1E 1740.7-2942”.

Method Flux The IBIS method 32 ± 4 mCrab The XIS method 31 ± 2 mCrab The FC method 28 ± 3 mCrab The 3 observations 32 ± 2 mCrab

98 Table 6.13. Summary of the ﬂux estimation of the known hard X-ray sources.

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

-1.8 -0.0 — — 501053010 C 18.5 — 1.52 ± 0.11 — -1.5 -0.0 — — 501052010 C 14.9 — 4.20 ± 0.14 — — — -0.9 -0.1 1E 1740.7-2942 — — 0.068 — 2.22 ± 0.28 -1.2 -0.0 — — 501051010 C 19.4 — 18.47 ± 0.18 — — — -0.4 -0.1 KS 1741-293 — — 0.015 — 0.08 ± 0.07 ± 99 — — -0.9 -0.1 1E 1740.7-2942 — — 0.469 — 15.29 1.91 -1.1 -0.0 — — 500019010 B 11.0 — 2.70 ± 0.15 — — — -0.9 -0.1 1E 1740.7-2942 — — 0.440 — 1.98 ± 1.80 -1.2 -0.3 — — 501057010 C 14.3 — 18.01 ± 0.19 — — — -0.4 -0.4 1A 1742-294 — — 0.022 — 0.33 ± 0.11 — — -0.9 -0.1 1E 1740.7-2942 — — 0.412 — 13.48 ± 1.72 -1.2 0.3 — — 501049010 C 16.2 — 4.13 ± 0.13 — — — -0.9 -0.1 1E 1740.7-2942 — — 0.063 — 2.05 ± 0.26 -0.8 0.3 — — 501048010 D 22.3 — 11.53 ± 0.14 — — — -0.9 -0.1 1E 1740.7-2942 — — 0.192 — 9.33 ± 0.78 -0.5 0.3 — — 501047010 D 17.8 — 9.37 ± 0.15 — — — -0.0 0.1 GRS 1741.9-2853 — — 0.019 — 0.27 ± 0.07 — — -0.4 -0.1 KS 1741-293 — — 0.192 — 0.86 ± 0.80 — — -0.9 -0.1 1E 1740.7-2942 — — 0.029 — 1.39 ± 0.12 -0.2 0.3 — — 501046010 D 22.9 — 14.90 ± 0.15 — Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — -0.1 -0.1 IGR J17456-2901 — — 0.214 — 0.56 ± 1.52 — — -0.0 0.1 GRS 1741.9-2853 — — 0.453 — 6.41 ± 1.83 — — -0.4 -0.1 KS 1741-293 — — 0.155 — 0.79 ± 0.64 -0.8 -0.3 — — 501056010 D 23.0 — 32.33 ± 0.20 — — — -0.4 -0.1 KS 1741-293 — — 0.072 — 0.30 ± 0.30 — — -0.4 -0.4 1A 1742-294 — — 0.159 — 2.07 ± 0.82

100 — — -0.7 -0.9 SLX 1744-299/300 — — 0.035 — 0.26 ± 0.14 — — -0.9 -0.1 1E 1740.7-2942 — — 0.500 — 25.17 ± 2.08 -0.8 -0.0 — — 501050010 C 17.5 — 27.05 ± 0.21 — — — -0.1 -0.1 IGR J17456-2901 — — 0.006 — 0.02 ± 0.04 — — -0.4 -0.1 KS 1741-293 — — 0.293 — 1.45 ± 1.28 — — -0.4 -0.4 1A 1742-294 — — 0.049 — 0.68 ± 0.24 — — -0.9 -0.1 1E 1740.7-2942 — — 0.613 — 19.95 ± 2.49 -0.6 -0.1 — — 500018010 B 43.8 — 17.58 ± 0.11 — — — -0.1 -0.1 IGR J17456-2901 — — 0.028 — 0.19 ± 0.19 — — -0.4 -0.1 KS 1741-293 — — 0.583 — 2.83 ± 2.47 — — -0.4 -0.4 1A 1742-294 — — 0.386 — 9.26 ± 1.96 — — -0.9 -0.1 1E 1740.7-2942 — — 0.443 — 2.13 ± 1.82 -0.5 -0.3 — — 501055010 D 19.9 — 34.59 ± 0.22 — — — -0.1 -0.1 IGR J17456-2901 — — 0.018 — 0.04 ± 0.14 Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — -0.4 -0.1 KS 1741-293 — — 0.427 — 1.18 ± 1.78 — — -0.4 -0.4 1A 1742-294 — — 0.725 — 9.33 ± 3.65 — — -0.7 -0.9 SLX 1744-299/300 — — 0.029 — 0.21 ± 0.11 — — -0.9 -0.1 1E 1740.7-2942 — — 0.275 — 13.68 ± 1.13 -0.2 -0.0 — — 100027020 A 33.4 — 18.09 ± 0.13 — — — 0.3 -0.0 1E 1743.1-2843 — — 0.180 — 1.03 ± 1.15

101 — — 0.2 -0.2 1A 1743-288 — — 0.168 — 0.18 ± 0.85 — — -0.1 -0.1 IGR J17456-2901 — — 0.684 — 3.59 ± 4.72 — — -0.0 0.1 GRS 1741.9-2853 — — 0.576 — 2.35 ± 2.27 — — -0.4 -0.1 KS 1741-293 — — 0.477 — 4.14 ± 2.01 — — -0.4 -0.4 1A 1742-294 — — 0.165 — 0.66 ± 0.80 -0.2 -0.0 — — 100037010 A 36.8 — 16.97 ± 0.12 — — — 0.3 -0.0 1E 1743.1-2843 — — 0.182 — 0.99 ± 1.16 — — 0.2 -0.2 1A 1743-288 — — 0.169 — 0.19 ± 0.85 — — -0.1 -0.1 IGR J17456-2901 — — 0.687 — 3.00 ± 4.68 — — -0.0 0.1 GRS 1741.9-2853 — — 0.578 — 1.26 ± 2.20 — — -0.4 -0.1 KS 1741-293 — — 0.475 — 3.82 ± 2.02 — — -0.4 -0.4 1A 1742-294 — — 0.164 — 0.71 ± 0.80 -0.2 -0.3 — — 501054010 D 21.5 — 25.14 ± 0.19 — — — 0.3 -0.0 1E 1743.1-2843 — — 0.014 — 0.07 ± 0.09 Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — 0.2 -0.2 1A 1743-288 — — 0.171 — 1.05 ± 0.97 — — -0.1 -0.1 IGR J17456-2901 — — 0.345 — 0.82 ± 2.55 — — -0.0 0.1 GRS 1741.9-2853 — — 0.111 — 1.66 ± 0.43 — — -0.4 -0.1 KS 1741-293 — — 0.371 — 1.24 ± 1.54 — — -0.4 -0.4 1A 1742-294 — — 0.516 — 6.59 ± 2.57 -0.2 -0.2 — — 501008010 C 102.6 — 19.32 ± 0.08 —

102 — — 0.3 -0.0 1E 1743.1-2843 — — 0.285 — 1.91 ± 1.97 — — 0.2 -0.2 1A 1743-288 — — 0.355 — 1.27 ± 1.90 — — -0.1 -0.1 IGR J17456-2901 — — 0.714 — 4.65 ± 4.90 — — -0.0 0.1 GRS 1741.9-2853 — — 0.405 — 0.72 ± 1.51 — — -0.4 -0.1 KS 1741-293 — — 0.312 — 1.50 ± 1.33 — — -0.4 -0.4 1A 1742-294 — — 0.219 — 4.29 ± 1.08 -0.1 0.2 — — 501009010 C 43.7 — 9.67 ± 0.10 — — — 0.3 -0.0 1E 1743.1-2843 — — 0.290 — 2.69 ± 2.13 — — 0.2 -0.2 1A 1743-288 — — 0.093 — 0.46 ± 0.52 — — -0.1 -0.1 IGR J17456-2901 — — 0.446 — 2.81 ± 3.15 — — -0.0 0.1 GRS 1741.9-2853 — — 0.770 — 0.88 ± 2.93 — — -0.4 -0.1 KS 1741-293 — — 0.072 — 0.25 ± 0.31 0.1 -0.1 — — 100027010 A 34.9 — 17.15 ± 0.13 — — — 0.6 -0.1 IGR J17475-2822 — — 0.135 — 0.27 ± 0.42 Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — 0.3 -0.0 1E 1743.1-2843 — — 0.676 — 3.08 ± 4.42 — — 0.2 -0.2 1A 1743-288 — — 0.587 — 0.92 ± 2.99 — — -0.1 -0.1 IGR J17456-2901 — — 0.607 — 2.95 ± 4.23 — — -0.0 0.1 GRS 1741.9-2853 — — 0.530 — 2.41 ± 2.14 — — -0.4 -0.1 KS 1741-293 — — 0.063 — 0.59 ± 0.26 0.1 -0.1 — — 100037040 A 33.0 — 17.06 ± 0.13 —

103 — — 0.6 -0.1 IGR J17475-2822 — — 0.134 — 0.41 ± 0.42 — — 0.3 -0.0 1E 1743.1-2843 — — 0.673 — 3.98 ± 4.30 — — 0.2 -0.2 1A 1743-288 — — 0.581 — 0.62 ± 2.94 — — -0.1 -0.1 IGR J17456-2901 — — 0.610 — 2.34 ± 4.10 — — -0.0 0.1 GRS 1741.9-2853 — — 0.535 — 0.90 ± 2.00 — — -0.4 -0.1 KS 1741-293 — — 0.063 — 0.45 ± 0.26 0.1 -0.1 — — 100048010 C 56.1 — 15.27 ± 0.10 — — — 0.6 -0.1 IGR J17475-2822 — — 0.140 — 0.45 ± 0.44 — — 0.3 -0.0 1E 1743.1-2843 — — 0.677 — 6.15 ± 4.69 — — 0.2 -0.2 1A 1743-288 — — 0.594 — 1.51 ± 3.12 — — -0.1 -0.1 IGR J17456-2901 — — 0.597 — 2.55 ± 4.09 — — -0.0 0.1 GRS 1741.9-2853 — — 0.536 — 0.67 ± 1.98 — — -0.4 -0.1 KS 1741-293 — — 0.074 — 0.48 ± 0.32 0.4 -0.1 — — 500005010 B 60.4 — 34.63 ± 0.13 — Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — 0.6 -0.1 IGR J17475-2822 — — 0.516 — 1.63 ± 1.64 — — 0.3 -0.0 1E 1743.1-2843 — — 0.625 — 0.23 ± 3.75 — — 0.2 -0.2 1A 1743-288 — — 0.623 — 5.85 ± 3.19 — — -0.1 -0.1 IGR J17456-2901 — — 0.183 — 1.48 ± 1.29 — — -0.0 0.1 GRS 1741.9-2853 — — 0.120 — 0.27 ± 0.45 — — 1.0 -0.5 IGR J17497-2821 — — 0.026 — 0.11 ± 0.11

104 0.6 -0.1 — — 100037060 A 65.9 — 8.49 ± 0.07 — — — 0.6 -0.1 IGR J17475-2822 — — 0.766 — 4.01 ± 2.39 — — 0.3 -0.0 1E 1743.1-2843 — — 0.193 — 0.98 ± 1.27 — — 0.2 -0.2 1A 1743-288 — — 0.068 — 1.10 ± 0.35 — — 1.0 -0.5 IGR J17497-2821 — — 0.144 — 0.37 ± 0.62 0.6 0.1 — — 501040010 C 49.8 — 7.01 ± 0.08 — — — 0.6 -0.1 IGR J17475-2822 — — 0.615 — 4.35 ± 2.06 — — 0.3 -0.0 1E 1743.1-2843 — — 0.198 — 0.64 ± 1.23 — — 0.2 -0.2 1A 1743-288 — — 0.008 — 0.03 ± 0.04 — — -0.0 0.1 GRS 1741.9-2853 — — 0.006 — 0.02 ± 0.02 0.8 -0.2 — — 501039010 D 84.6 — 6.25 ± 0.06 — — — 0.6 -0.1 IGR J17475-2822 — — 0.580 — 2.90 ± 1.99 — — 0.3 -0.0 1E 1743.1-2843 — — 0.127 — 0.28 ± 0.81 — — 0.2 -0.2 1A 1743-288 — — 0.070 — 0.21 ± 0.37 Table 6.13—Continued

Position Sequence number Term Expb Angular Detected Point Observaiona Source or [ks] response ﬂux Sources l(◦) b(◦) l(◦)(b◦) Source name ratio [mCrab] [mCrab]

— — 1.0 -0.5 IGR J17497-2821 — — 0.331 — 1.63 ± 1.45 1.0 -0.1 — — 100037070 A 8.6 — 4.18 ± 0.18 — — — 1.6 0.1 AX J1749.1-2733 — — 0.052 — 0.09 ± 0.11 — — 0.6 -0.1 IGR J17475-2822 — — 0.172 — 0.73 ± 0.54 — — 1.0 -0.5 IGR J17497-2821 — — 0.245 — 0.55 ± 1.05 1.2 0.0 — — 501059010 D 49.6 — 3.55 ± 0.07 —

105 — — 1.7 0.1 AX J1749.2-2725 — — 0.043 — 0.08 ± 0.09 — — 1.6 0.1 AX J1749.1-2733 — — 0.133 — 0.32 ± 0.29 — — 0.6 -0.1 IGR J17475-2822 — — 0.126 — 0.59 ± 0.40 — — 1.0 -0.5 IGR J17497-2821 — — 0.150 — 0.70 ± 0.67 1.3 0.2 — — 501058010 D 46.8 — 2.50 ± 0.07 — — — 1.7 0.1 AX J1749.2-2725 — — 0.138 — 0.27 ± 0.29 — — 1.6 0.1 AX J1749.1-2733 — — 0.271 — 0.64 ± 0.59 1.5 0.0 — — 501060010 D 50.2 — 2.15 ± 0.07 — — — 1.7 0.1 AX J1749.2-2725 — — 0.395 — 0.70 ± 0.82 — — 1.6 0.1 AX J1749.1-2733 — — 0.621 — 1.31 ± 1.35

aFOV (XIS-nominal) center position.

bEﬀective exposure of the HXD-PIN data. Class B

Class A

AX J1749.2-2725 GRS 1741.9-2853 1E 1743.2843 AX J1749.1-2733 IGR J17475-2822 KS 1741-293 1E 1740.7-2942 1A 1742-288 1A 1742-294 IGR J17497-2821

SLX 1744-299/300

Figure 6.37: The positions of the hard X-ray sources and the HXD-PIN FOVs of the all the observations. 6.4 Analysis of the Hard X-ray diﬀuse emission

In this section, we classify the observations into three classes based on the number of IBIS catalogued hard X-ray point sources (section 6.3.1) in the HXD-PIN FOV. In order to estimate a systematic error on the flux contamination by point sources, we analyze the data for each class independently (Figure 6.37). Firstly, we select the observation which has no source in the FOV and define as class A. Only one observation is classified into class A (”501053010”) and lies in (l, b) = (−1.8◦, 0.0◦). Secondly, we define the observation in the west region which have one contaminated source with angular-response-corrected flux above 1 mCrab as class B. Since there are multiple hard X-ray point sources especially in the center and east region, the systematic uncertainties could be large and comparable with the detected fluxes. Therefore, we define these observations as class C. We show the analyses in the order of class A, B and C.

6.4.1 Class A The detected flux of the class A observation is considered as the hard X-ray diffuse emission if the detection is significant compared with statical and systematic uncertainties. The detected flux is 1.5 ± 0.1 mCrab in the energy range of 14 – 40 keV (Table 6.6). And the corresponding surface brightness is 6.8 ±0.5 × 10−11 erg/sec/deg2. The flux is detected with 15 σ significance level and 5 times as large as the systematic uncertainty of the NXB modeling. Hence, the detection in 14 – 40 keV is significant. We verify if the lightcurve fluctuates significantly (Figure 6.16). one σ of the fluctu-

106 Table 6.14. Summary of the known hard X-ray sources in the FOV for the west region.

Class Position Sequence Terma Source Detectedb Elong.c Ratiod Point source e Diﬀuse f Occup.g l(◦) b(◦) number name [mCrab] [mCrab] [mCrab]

A -1.8 -0.0 501053010 C Nothing 1.5 ± 0.1 — 0.0 — 1.5 ± 0.1 0% B -1.5 0.0 501052010 C 1E 1740 4.2 ± 0.1 37.2′ 0.07 2.1 ± 0.1 2.1 ± 0.2 50% -1.2 0.0 501051010 C 1E 1740 18.5 ± 0.2 17.8′ 0.47 15.0 ± 1.0 3.5 ± 1.0 81% -1.1 -0.0 500019010 B 1E 1740 2.7 ± 0.2 14.2′ 0.44 0.0 2.7 ± 0.2 0% -1.2 0.3 501049010 C 1E 1740 4.1 ± 0.1 31.3′ 0.06 2.0 ± 0.1 2.1 ± 0.2 49% -1.2 -0.3 501057010 C 1E 1740 17.7 ± 0.2 17.8′ 0.41 13.2 ± 0.8 4.9 ± 0.8 75% ′ 107 -0.8 0.3 501048010 D 1E 1740 11.5 ± 0.1 13.8 0.19 9.0 ± 0.8 2.5 ± 0.8 78% -0.5 0.3 501047010 D 1E 1740 9.4 ± 0.2 33.4′ 0.03 1.4 ± 0.1 8.0 ± 0.2 15% -0.2 0.3 501046010 D GRS 1741 14.9 ± 0.2 14.9′ 0.45 6.4 ± 1.8 8.5 ± 1.8 43%

aB; 2006/02 – 2006/03. C; 2006/09 – 2006/10. D; 2007/03.

bFluxes subtracted NXB and CXB. cElongations between the positions of the observations and the contaminated point source.

dRatios of the angular response for the contaminated point source. eEstimated contamination from the contaminated point source.

f Residual fluxes subtracted the contaminations. gRates occupied by the contaminations in the detected fluxes. Figure 6.38: The HXD-PIN spectra of the class A observation ”501053010”. ation in the lightcurve is 0.4 mCrab. While, the statical uncertainties of the individual time bins (5.7 ksec) is ∼ 0.3 mCrab. Considering the systematic uncertainty of the NXB modeling of 0.3 mCrab (Mizuno et al. 2006), the fluctuation seen in the light curve is comparable of the uncertainty of our NXB prediction. Since the systematic uncertainty is estimated by the integrated fluxes in 10 ksec (Mizuno et al. 2006), the estimation with time bin size of 5.7 ksec could be conservative. From these consideration we conclude that the detected flux is not influenced by unknown transient hard X-ray sources with hour scale fluctuation. We analyze the HXD-PIN spectrum. Figure 6.38 shows the spectra of raw fluxes, the NXB and CXB fluxes, detected fluxes (raw fluxes minus the NXB and CXB fluxes) and the 3 % level of the NXB. The detected fluxes extend to around 30 keV, and 3 times larger than the systematic uncertainties of the NXB, the 3% level. Therefore, in this spectrum, the statical error is dominant and the fluxes are detected with at least 3 σ significance up to 30 keV. The photon index of the spectrum fitted with a power-law function is 2.3 ± 0.5. If we adopt a thermal bremsstrahlung model, the resultant temperature calculated +18 to be 20−7 keV. We analyze the wide band spectra of XIS and HXD-PIN. As described in section 6.1.4, the difference of the FOVs between XIS and HXD-PIN is adjusted by the constant factor calculated from the XIS surface brightness distribution. The constant factor of this observation is calculated to be 0.96, which is consistent with the assumption that the emission in the HXD-PIN FOV is almost uniform. Figure 6.39 shows the wide band spectra compared with the hot plasma model plus the neutral lines without the power-law component. Obviously, there is an excess toward the thermal emission. We then perform the spectral fitting to the wide band spectra with ”Model A” which consists of the hot plasma model, the neutral lines and the power-law component. We

108 a . e n . oa epcieywihaetevle ftecne ein(oaaet 0.96. (Koyama be region to center fixed the is of be factor values to constant the abundance the are this iron And which In and respectively 2007c). 6.15. plasma al. solar Table hot in 1.2 the listed and of are keV temperature parameters 6.5 the fit assume best we The 2.3 fitting, level. spectral becomes be confidence index can % 1.2 90 photon spectra becomes with The The keV, component power-law function. the 6.40. power-law of Figure a flux by in reduced fitted A a well Model with be A with Model fitted by spectra fitted band wide the show emCi siae ihfu ehd n h eutn ue r ossetamong consistent are in fluxes 1740.7-2943” resultant ”1E the of and flux the methods 6.3.4, four section with in estimated 1740.7-2943” described is ”1E As is C source 6.41. contaminated term the Figure B, in class as shown defined as observations 8 of out 7 In B Class 6.4.2 10 is unit The al .5 h etfi aaeeso h iebn pcrlfitn o the for fitting spectral band wide the observation. of A parameters class best-fit The 6.15: Table ”501053010” observation A class model. the plasma of hot the spectra with band compared wide The 6.39: Figure . fie)12(xd 2.3 (fixed) 1.2 (fixed) 6.5 T[e]Z [keV] kT − 10 erg/cm

χ normalized counts/sec/keV −2 0 2 10−4 10−3 0.01 0.1 2 /sec e F ± χ Flux Γ 1 0 2 0.1 f4./3=07.Teecs opnn could component excess The 0.70. = 44.2/63 of c h a n 109 n e l

e n − +0 e r 0 g 5 . . 0.3 y 1 2 .

5 ( k − Ca nteeeg ag f1 40 – 14 of range energy the in mCrab e V 11keV ) 2 0 a Flux . .0(63) 0.70 0.2 14 − 40keV a ± χ ν 2 .,adthe and 0.1, ( ν ) htnidcsaelse nTbe61.I hstbe h oiatucranyi statical is uncertainty resultant dominant the The table, this spectrum. In observational 6.16. each Table in the listed for are indices performed photon is function power-law a decrease subsection. to next the tend the adding in fluxes of analyzed C These is distribution class distribution brightness brightness 6.45. of surface surface data Figure X-ray observational The in hard The shown . center. is B the 3 emission class leaving above diffuse as with X-ray defined detected hard observations are the the emission all diffuse in X-ray 1999) level hard the al. as et considered ”GRS (Cocchi of work index previous photon the 6.8. also The Table spectrum from in et the 6.3.4). 1.6 listed from Santo (section being as value as (Del method the assumed work with XIS being is consistent previous the 1741.9-2853” is as with the index assumed 1740.7-2942” from photon is The ”1E correlation 1740.7-2942” of 6.8). flux (Table ”1E and 2005) of index source index al. photon point photon the the of The from shape also spectral 1.5 function. are the power-law estimations assume These we a addition, adopted. as In is 6.14. method ”GRS Table and IBIS in 1740.7-2942” consistent the summarized ”1E D”, are of ”term fluxes other the during residual each for 1741.9-2853” And, as the negligible. IBIS same and is with the ”500019010” ”500019010” observation detected almost of not are flux is 1.2 positions detected flux with whose the the ”501051010” during While, Since of 1740.7-2942” The method. ”1E flux significantly. IBIS 6.42. of B the flux Figure term with The in during estimated shown ratio 6.14. is as Table the D observations by in term flux individual listed the the are of contributions for product calculated response from calculated angular are the contributions of The methods. the all eaayetesetao h adXrydffs msin h pcrlfitn with fitting spectral The emission. diffuse X-ray hard the of spectra the analyze We are which fluxes residual the contaminations, source the of subtraction the After iue64:Tewd adsetao h ls bevto ”501053010” observation A class the of A. Model spectra with band fitted wide The 6.40: Figure σ infiac ecnie httecnrbto rm”E70724”t the to ”1E1740.7-2942” from contribution the that consider we , significance

χ normalized counts/sec/keV −2 0 120−4 10−3 0.01 0.1 1 0 c h a n 110 n e l

e n e r g y

( k e V ) 2 0 σ signiﬁcance Contribution＝1mCrab ( Obs. term： 2006/02、2006/10、2007/03)

HXD-PIN FOV of (-1.1, 0 ) (-0.2, +0.3) (-0.5, +0.3) (-0.8, +0.3) (-1.2, +0.3)

GRS 1741 (-1.2, 0) Class A (-1.5, 0) (-1.8, 0) (-1.1, 0) (0, 0) KS 1741 1E 1740.7-2942

(-1.2, -0.3) 1A 1742

Figure 6.41: Perspective of the contributions from the hard X-ray point sources in class B observations.

40 1

30 401048010 401047010 0.8 20 401046010 10 501057010 500019010 0.6 (’)

Y - 0 Y K

S 0.4 -10 501051010

-20 501052010 0.2 -30 501049010

-40 0 -40 -30 -20 -10 0 10 20 30 40 SKY-X (’)

Figure 6.42: Two-dim proﬁle of HXD-PIN angular response and the position of the hard X-ray sources in the observations of class B.

111 Figure 6.43: Two-dim profile of HXD-PIN angular response and the position of the hard X-ray sources in the observations of class B. error except for ”501046010”. We analyze the individual wide band spectra of the Galactic diffuse emission like class A. The temperature of the hot plasma and iron abundance are assumed to be 6.5 keV and 1.2 solar respectively again. The constant factors of the individual observations are calculated from the soft X-ray surface brightness distribution, and listed in Table 6.17. Five of eight spectra can be fitted with model A with the calculated constant factors. Figure 6.44 shows the 5 wide band spectra. The photon indices of the HXD- PIN spectra of these 5 observations are dominated by the statical uncertainties, not by systematic uncertainties of the NXB or the point source subtraction, as listed in Table 6.6. Hence, we consider only the statical errors in these 5 wide band spectra. The resultant parameters of the power-law component are listed in Table 6.17. The photon indices and the fluxes of the power-low component are determined precisely like the class A spectra. The individual photon indices of the wide band spectra are shown in Figure 6.17. The photon indices are between 1.8 – 2.5 at 90 % confidence error. And, any hard X-ray spectra do not show a sign of cut-off at least below 30 keV. In addition, we have obtained the flux of the power-law component and the ratio of thermal : power-law = 1:4 in the energy range of 14 – 40 keV. While, the wide band spectra of the other 3 observations can not fitted with Model A utilizing the calculated constant factors. When the constant factors are thawed, the wide band spectra can be fitted with Model A (Figure 6.46). The constant factors become ∼ 2 as listed in Table 6.17. The positions of the 3 observations depart from the Galactic plane. Since the soft X-ray surface brightness distribution along the Galactic latitude is determined by the 3 data points as shown in Figure 6.17, it is considered that systematic error of the distribution is large.

112 Table 6.16: Summary of power-law model ﬁtting to the hard X-ray diﬀuse emission.

a 2 Class Position Sequence Γ Systematic χν (ν) l(◦) b(◦) number NXBb Point sourcec A -1.8 -0.0 501053010 2.3 ± 0.5 ± 0.1 — 0.79 (26) B -1.5 -0.0 501052010 2.0 ± 0.3 ± 0.1 ± 0.1 0.95 (27) ± ± +0.3 -1.2 -0.0 501051010 2.5 0.2 0.1 −0.2 1.43 (45) +0.5 ± -1.1 -0.0 500019010 2.7 −0.4 0.1 — 1.54 (19) ± ± +0.3 -1.2 -0.3 501057010 3.1 0.3 0.1 −0.2 1.09 (39) -1.2 0.3 501049010 2.5 ± 0.3 ± 0.1 ± 0.1 1.17 (29) +0.3 ± +0.3 -0.8 0.3 501048010 2.2 −0.2 0.1 −0.2 0.80 (44) -0.5 0.3 501047010 1.9 ± 0.1 ± 0.0 ± 0.1 1.38 (37) ± ± +0.4 -0.2 0.3 501046010 2.9 0.1 0.0 −0.2 0.59 (44)

aErrors represent 90% conﬁdence. bSystematic errors for ± 3% of the NXB level. cSystematic errors for the contaminations from the point sources.

Table 6.17: The resultant parameters from the ﬁtting of the wide band spectra.

a b 2 Class Position Sequence Spectral Const. Calc. Γ Flux χν (ν) l(◦) b(◦) number ID A -1.8 -0.0 501053010 1 0.96 (fix) 0.96 2.3 ± 0.2 1.2 ± 0.2 0.70 (63) B -1.5 -0.0 501052010 2 0.92 (fix) 0.92 2.0 ± 0.2 1.8 ± 0.3 1.16 (64) -1.2 -0.0 501051010 3 0.85 (fix) 0.85 1.9 ± 0.1 1.8 ± 0.3 1.24 (101) -1.1 -0.0 500019010 4 0.88 (fix) 0.88 2.3 ± 0.2 1.9 ± 0.3 1.32 (50) -1.2 -0.3 501057010 5 2.7 ± 0.6 1.06 2.4 ± 0.3 3.4 ± 1.4 1.36 (80) -1.2 0.3 501049010 6 1.07 (fix) 1.07 2.0 ± 0.2 1.7 ± 0.3 1.42 (60) -0.8 0.3 501048010 7 1.14 (fix) 1.14 2.0 ± 0.2 2.1 ± 0.3 0.82 (79) -0.5 0.3 501047010 8 2.0 ± 0.3 1.33 1.7 ± 0.1 6.8 ± 1.1 1.54 (73) -0.2 0.3 501046010 9 3.4 ± 0.4 1.35 2.5 ± 0.1 5.9 ± 1.5 0.77 (86)

aCalculated constant factors. bFluxes of the power-law component.

113 χ normalized counts/sec/keV χ normalized counts/sec/keV iue64:Tewd adsetao h ls bevto hc a be can which observation B class factors. constant the calculated of the spectra with band ﬁtted wide The 6.44: Figure −2 0 2 10−4 10−3 0.01 0.1 −2 0 2 10−4 10−3 0.01 0.1 1 0 1 0 c c h h a a n n n n e e l l

e e n χ normalized counts/sec/keV n e e r r g −2 0 2 10−4 10−3 0.01 0.1 g y y

( ( k k e e V V ) ) 2 2 0 0 5 5 0 0 1 1 0 0 0 0 5 1 2 9 0 0 c 1 1 h 0 0 a 114 n n e l

n χ normalized counts/sec/keV χ normalized counts/sec/keV e r g −2 0 2 104−4 10−3 0.01 0.1 −2 0 2 140−4 10−3 0.01 0.1 y

( k e V ) 2 0 5 0 1 1 1 0 0 0 4 8 0 1 0 c c h h a a n n n n e e l l

e e n n e e r r g g y y

( ( k k e e V V ) ) 2 2 0 0 5 5 0 0 1 1 0 0 4 5 9 1 0 0 1 1 0 0 Figure 6.45: The resultant photon indices from the wide band spectral fitting. The spectral IDs correspond to the number in Table 6.17. The filled data points mean that the wide band spectra can be fitted with the calculated constant factors. The outline data points mean that the wide band spectra can not be fitted with the calculated constant factors. 1 1 . . 0 0 V V e e k k / / c c e e 1 1 s s / / 0 0 . . s s t t 0 0 n n u u o o c c d d e e 3 3 z z − − i i l l 0 0 a a 1 1 m m r r o o n n 4 4 − − 0 0 4 1 1 2 2 0 χ χ 0 2 − 2 −

10 20 10 20 channel energy (keV) channel energy (keV) 1 . 0 V e k / c e 1 s / 0 . s t 0 n u o c

d e 3 z − i l 0 a 1 m r o n 4 − 0 1 2 χ 0 2 −

10 20 channel energy (keV)

Figure 6.46: The wide band spectra of the class B observation which can not be ﬁtted with the calculated constant factors.

115 Figure 6.47: The surface brightness of the hard X-ray emission. The upper panel shows the distribution of the detected fluxes (filled circle) and the fluxes of the estimated contaminations of the known hard X-ray sources (outline circle). The lower panel shows the distribution of the residual fluxes, the hard X-ray diffuse emission. The data points of red and black markers mean |b| < 0.2◦ and |b| > 0.2◦ respectively.

6.4.3 Class C

The surface brightness distribution of the hard X-ray diffuse emission are obtained to extend the analysis to class C observations. We estimate the flux due to source contaminations in the class C observations based on the IBIS method. The surface brightness distribution of the residual fluxes, the hard X-ray diffuse emission, is shown in Figure 6.47. Here we use conservative systematic errors of ∼ 4 mCrab for the individual sources. As result, the residual fluxes have large errors in the center region and east region where there are multiple hard X-ray sources.

6.4.4 Distribution of Hard X-ray Emission

We show the soft and hard X-ray surface brightness distribution in Figure 6.48 and Figure 6.49. These ﬁgures show the soft X-ray surface brightness distribution and the hard X-ray one compared with the distribution model rounded by the individual FOVs as described in section 6.1.4. The soft X-ray distribution is consistent with the distribution model within ﬁtting error of 20%. While, the errors of the hard X-ray distribution

116 Figure 6.48: Top panel; the soft X-ray surface brightness distribution (red data points) compared with the distribution rounded by the XIS FOV (blue line). The blue dashed line means the fitting errors of 20 %. Bottom panel; the ratios of the red data points to the blue line. is too large to be compared with the model quantitatively. The errors of the sources fluxes with the IBIS method consist of the fitting errors and the uncertainties due to time variability with short time scale in the IBIS lightcurves (section 6.3.2). This short- fluctuation uncertainties mean the uncertainties of flux fluctuations within the interval of the IBIS lightcurve (every 1–4 days). It is difficult to consider that all the hard X-ray sources in the HXD-PIN FOV becomes bright at the same time and fluctuate at the time scale of a few days. Our estimation of systematic error could be too conservative. However, even if we assume the hard X-ray sources do not fluctuate within a few days and evaluate the hard X-ray distribution with only the fitting errors (Figure 6.50) for the data points, the systematic errors are still large and we could not draw strong statement on the the hard X-ray distribution. If we use the data points composed of class A and class B, it is likely that soft X-ray and hard X-ray follow the similar distribution. Assuming the hard X-ray distribution is same as that of soft X-ray, the total luminosity of the hard X-ray is calculated to be 5 ±0.5×1036 erg/sec in the energy range of 14 – 40 keV and the region of |l| < 2.0◦, |b| < 0.5◦. And, assuming the mixing ratio of thermal and power-law component is 1:4 in the whole region, the total luminosity of the power-law component is calculated to be 4 ± 0.4 × 1036 erg/sec. As a further estimation, we try to compare the hard X-ray surface brightness distribution with the soft X-ray one quantitatively. As described in section 6.1.4, we use a form of G × exp(−|l|/l0) × exp(|b|/b0) + C to model the surface brightness distribution of the XIS fluxes. In order to model the hard X-ray distribution, we start from the XIS flux

117 Figure 6.49: Top panel; the hard X-ray surface brightness distribution (red data points) compared with the soft X-ray one rounded by the HXD-PIN FOV (blue line). The blue dashed line means the ﬁtting errors of 20 %. Bottom panel; the ratios of the red data points to the blue line.

◦ ◦ distribution (G = 9.3, C = 0.8, l0 = 0.33 and b0 = 0.15 ). However, with the HXD-PIN data, it is difficult to constrain the scale length (l0) and the scale height (b0). Therefore, ◦ ◦ we use the parameters from the XIS (l0 = 0.33 and b0 = 0.15 ). The constant C is calculated from the assumption that the emission is almost uniform, and therefore, we use C = 0.8. Here, we assume the function as g × G × exp(−|l|/l0) × exp(|b|/b0) + C, and we compare g with 1.0 from the soft X-ray distribution∑ after spatial fitting for the hard X-ray distribution. In the spatial fitting, we define P as {(data − model)/error}2, and search the value of g which gives minimum of P . We should note that the distribution P dose not follow the χ2 distribution because of overestimation of systematic error. Figure 6.51 shows the P distribution for g and P becomes minimum at g = 0.7. If we adopt g as the optimum value, it is suggested that the X-ray spectra could be softer around the center. In the future, that could be shown by further hard X-ray observations with small systematic errors.

118 Figure 6.50: The hard X-ray surface brightness distribution in the case that the IBIS ﬂuxes have only the ﬁtting errors not systematic those. The green line shows the rounded distribution assuming the real distribution as g = 0.7.

Figure 6.51: The P distribution for the factor of the distribution model.

119 Chapter 7

Discussion

7.1 Brief Summary of the Observational Results

We analyzed the 29 observational data collected from the Galactic center observation project with the Suzaku team. The following results are obtained in the present analysis.

• We detect the hard X-ray diﬀuse emission in the Galactic center region of |l| < 2.0◦ and |b| < 0.5◦. The intensity is calculated to be 1.5 – 10 mCrab which corresponds to 0.7 – 4.5×10−10 erg/cm2/sec/deg2.

• The overall spectrum is well ﬁtted with by the model composed of a hot plasma, neutral lines and a power-law function. The photon index obtained from each position ranges from 1.8 to 2.5 and the mixing ratio of thermal and power-law component is 1 : 4 in the energy range of 14 – 40keV. Any hard X-ray spectra do not show a sign of cut-oﬀ at least below 30 keV.

• The hard X-ray surface brightness distribution is obtained near the Galactic center. Although we could not draw strong statement on the the hard X-ray distribution because of large systematic errors, it is likely that soft X-ray and hard X-ray follow the similar distribution. The total luminosity of the power-law component is calculated to be 4 ± 0.4 ×1036 erg/sec (14 – 40 keV, |l| < 2.0◦, |b| < 0.5◦).

7.2 Uncertainties of the Hard X-ray Diﬀuse Emis- sion

In order to verify the signiﬁcance of the hard X-ray emission, we estimate uncertainties of the detected hard X-ray diﬀuse emission by class A, B, C.

120 Class A In the detected flux of class A, there are three possible uncertainties as listed in Table 7.1. The first one is the reproducibility of the NXB model (0.3 mCrab, Mizuno et al. (2006)). This uncertainty is too small to explain the detected flux of 1.5 ± 0.1 mCrab. The second one is a contribution from transient hard X-ray sources which are not detected monitoring observations with IBIS. In addition to the actual source list, the lightcurve of the observation is used to confirm the variability of sources (section 6.4.1). We concluded that the transient source, if exists, should be dimmer than the IBIS detection limit of 4 mCrab which corresponds to ∼ 2 mCrab if the source locates outside the FOV of HXD-PIN where the ratio of the angular response is above ∼ 0.5. The third one is the superposition of dim point sources. As described in section 7.3, the estimated contribution is ∼ 0.3 mCrab and could not explain the detected flux. This possibility is discussed in the next section closely.

Table 7.1: Uncertainties of class A

Detected ﬂux 1.5 ± 0.1 mCrab Reproducibility of the NXB model 0.3 mCrab Not-cataloged transient hard X-ray source Diﬃcult Dim point sources (< 1 mCrab) ∼ 0.3 mCrab

Class B In the class B observations, a dominant uncertainty is the flux estimation of the hard X-ray point source. There are two uncertainties as listed in Table 7.2. The first one is the systematic uncertainty of the angular response (∼ 1%). Since the fluxes of the hard X-ray sources (”1E1740.7-2942” or ”GRS 1741.9-2853”) are estimated to be 30 – 50 mCrab, the contribution is 0.5 mCrab at maximum. The uncertainty could not explain the detected fluxes in the class B observations.

Table 7.2: Uncertainties of class B

Detected ﬂux 2 – 8 mCrab Systematic uncertainty of the angular response (1%) < 0.5 mCrab Flux estimation of the hard X-ray sources 0.1 – 1.0 mCrab

The second one is the uncertainties of the flux estimation of the hard X-ray sources. These uncertainties are different for each observation, but within 0.1 – 1.0 mCrab. Hence, these are much smaller than the diffuse flux estimated in class B data.

121 Summarized, it is difficult to deny the detection of the hard X-ray diffuse emission also in the class B observations. Hence, the hard X-ray diffuse emission exists certainly in the west region widely.

Class C In the class C observations, the uncertainties of the IBIS lightcurves are dominant. There are two uncertainties; the fitting errors and the short fluctuation uncertainties as listed in Table 7.3. When we consider the short fluctuation, errors are comparable to the detected fluxes as shown in Figure 6.47 and 6.49. If we consider that the short fluctuation uncertainties are too conservative and adopt only statical errors, the detection of the hard X-ray diffuse emission of the class C observation becomes significant as shown in Figure 6.50.

Table 7.3: Uncertainties of class C (including possible systematic errors)

Detected ﬂux 2 – 10 mCrab Fitting errors 2 – 5 mCrab Short ﬂuctuation uncertainties 4 – 10 mCrab

7.3 Contributions from the Dim Point Sources

The most possible celestial class of objects that could explain the detected diffuse hard X-ray in the Galactic center region is Cataclysmic Variable (CV). The CV is classified into three categories by its magnetic field strength; non-magnetic CV with < 105 Gauss, Intermediate Polar (IP) with 105−7 Gauss and Polar with > 107 Gauss (Talbe 7.4, Ezuka & Ishida (1999)). The continuum component in CV spectra is characterized to be a thermal bremsstrahlung. The temperatures of non-magnectic, IP and Polar are < 10 keV, 10 – 40 keV and 10 – 40 keV respectively. Since the spectral shape of the hard X-ray diffuse emission is above 10 keV (section 6.4.1), the possibility of non-magnetic CVs is denied. In addition, discovered Polars are much less than IPs, and therefore, only IPs are considered as the candidate of dim point sources. The total luminosity detected hard X-ray emission in the Galactic center region (5 ×1036 erg/sec in |l| < 2.0◦ and |b| < 0.5◦); which corresponds to 5000 IPs with a luminosity of 1033 erg/sec. Considering the energy band of 14 – 40 keV, the total luminosity becomes ∼ 3 times larger than 5 ×1036 erg/sec and the corresponding number of IPs becomes 15000. While, the number of the detected X-ray point sources with Chandra is 1352 in the region of |l| < 1.0◦ and |b| < 0.4◦ (Muno et al. 2006). Muno et al. (2006) mention that 90% of the total point sources are CVs. Therefore, the detected

122 number of CVs is calculated to be ∼ 1200. And, the detection limit of the observations is ∼ 1 × 1032 erg/sec. Hence, it is possible that almost all the detected CVs are composed of IPs and almost all IPs in the Galactic center region are already detected. In order to adjust the diﬀerence of the observational areas between Suzaku and Chandra, that is ”|l| < 2.0◦, |b| < 0.5◦” and ”|l| < 1.0◦, |b| < 0.4◦”, we utilize the number density around (l, b) = (0.5◦, 0◦) of 140 /deg2 (Muno et al. 2006). The number corresponding to the diﬀerence of the size between two area is calculated to be 140 /deg2 × 2.4 deg2 = 340. Finally, the number of 1540 (1200 + 340) is compared with 5000 – 15000 corresponding to the detected luminosity, and the superposition of the dim point sources could explain only 10 – 30 % of the observational luminosity.

Table 7.4: Kind of CVs (Ezuka & Ishida 1999)

Kind Magnetic ﬁeld Temperature Luminosity Non-magnetic < 105 Gauss < 10 keV 1031 erg/sec IP 105−7 Gauss 10 – 40 keV 1032−33 erg/sec Polar > 107 Gauss 10 – 40 keV 1032−33 erg/sec

7.4 Interpretation of the Hard X-ray Diﬀuse Emis- sion

There is no conclusive theory to explain the hard X-ray diﬀuse emission in the Galactic center region. However, we suggest that the thermal emission and non-thermal emission are correlated with each other because of the resemblance of the surface brightness distribution (section 6.4.3), and the two entirely diﬀerent emissions could be caused by the same origin. One possibility is that a physical phenomenon which heats the hot plasma and accelerate particles to high energy at the same time exists in entire the Galactic center region. Another possibility is that the emission is explained by a superposition of unknown sources which are too dim to be resolved spatially. We show the following restrictions from the observational results.

• The Galactic diﬀuse X-ray emission could be explained by highly ionized lines, a thermal bremsstrahlung with temperature of ∼ 6.5 keV and a power-law component.

• In the power-law component, there is no cut-oﬀ at least up to 30 keV.

• The photon index of the power-law component is between 1.8 – 2.5 with typical statical error of ± 0.2 and systematic error of ± 0.2.

123 • Total luminosity of the power-law component is 4 ±0.4×1036 erg/sec (14 – 40 keV, |l| < 2.0◦, |b| < 0.5◦).

124 Chapter 8

Conclusion

We performed the hard X-ray observations of the Galactic center region with HXD- PIN onboard Suzaku extending to |l| < 2.0◦ and |b| < 0.5◦. After subtraction of the contributions from hard X-ray point sources, we detected significant hard X-ray fluxes in the whole observed region. According to the wide band spectral analysis of XIS and HXD- PIN, it is shown that the power-law component exists and the derived photon indices are between 1.8 – 2.5. In the power-law component, there is no cut-off at least up to 30 keV. Although we could not draw strong statement on the the hard X-ray distribution because of large systematic errors, it is likely that soft X-ray and hard X-ray follow the similar distribution. The total luminosity of the power-law component is calculated to be 4 ± 0.4 ×1036 erg/sec (14 – 40 keV, |l| < 2.0◦, |b| < 0.5◦). In order to constrain the contributions of point sources to the diffuse emission, we study the Cataclysmic Variables (CVs) as a possible class of candidate. According to the analysis for the shape of the detected hard X-ray spectra, the kind of the CV is limited to Intermediate Polar (IP). Based on the recent results with Chandra, the luminosity integrated for all IPs in the Galactic center region could amount to only 10 – 30% of the detected hard X-ray luminosity. Although there is no conclusive theory to explain the hard X-ray diffuse emission in the Galactic center region, observational results derived in this thesis suggest two possibilities; (1) a physical phenomenon which heats the hot plasma and accelerate particles to high energy simultaneously exists in the Galactic center region, (2) is caused universally in unknown class of point sources which are too dim to be resolved spatially by current X-ray satellites. In order to solve the origin of the hard X-ray diffuse emission, it is necessary to reduce the systematic uncertainties of the hard X-ray point sources and obtain detailed hard X-ray surface brightness distribution. As the future work, observations with a hard X-ray imager of light-focused type which can resolve point sources and detect extended emissions at the same time are promising.

125 Appendix A

Image of Galactic Center Region with Swift

We have created a hard X-ray image of the Galactic center region with BAT onboard Swift. We have utilized archival data set in the web site of the Goddard Spcace Flight Center (http://heasarc.gsfc.nasa.gov/docs/swift/archive/). We have summed the archival data with ∼ 3 Msec and made the image as shown in Figure A.1.

126 0 . 6 0 0 0 0 4 . 0 0 0 2 GRS_17. 41.9_2853 5 0 1_2733 SAX_J1744.7_2916 1E_1743.1_2843 1.500 1.000 IGR_J17475_02.580202 IGR0_.0J0107456_2901 359.500 359.000 358.50 KS_1741_293 1E_1740.7_2942

XTEJ1748_288 0 SAX_J1747.0_2853 0 2 . 0 -

1A_1742_294 0 0 4 . 0 - - 0 . 6 0 0 - 0 . 8 0 0 SLX_1744_300 IGR_J17507_2856 0 0 0 .

1 IGRJ1746 - 0 0 2 . 1 -

1 2 3 4 5 6 7 8 9 10 11

Figure A.1: Image of the BAT

127 Appendix B

Individual Observational Data

B.1 The XIS spectra and the IBIS ﬂuxes of the bright point sources

128 χ normalized counts/sec/keV χ normalized counts/sec/keV χ normalized counts/sec/keV χ normalized counts/sec/keV −4 −2 0 2 0.05 0.1 0.2 0.5 −2 −1 0 1 5×10−3 0.01 0.02 0.05 0.1 −4 −2 0 2 4 0.01 0.1 1 −2 0 210−4 10−3 0.01 0.1 5 iueB1 h I pcr ftekonhr -a sources. X-ray hard known the of spectra XIS The B.1: Figure 5 5 5 1 0 c c c c h h h h a a a a 1 n n n n 1 0 1 n n n n 0 0 e e e e l l l l

e e e e

χ normalized counts/sec/keV n n n n e e e e r r r r −4 −2 0 2 40.01 0.1 1 g g g g y y y y

( ( ( ( k k k k e e e e 2 V V V V 0 1E 1743.1-2843(A1) ) ) ) ) 2 2 1E 1740.7-2942(C) 2 0 0 KS 1741-293(A2) 0 1A1742-294 (A1) 5 5 5 5 5 0 c 0 1 0 0 h 0 a 129 n n e l

n χ normalized counts/sec/keV χ normalized counts/sec/keV χ normalized counts/sec/keV χ normalized counts/sec/keV e r g −2 0 2 4 10−3 2×10−3 5×10−3 0.01 −4 −2 0 2 4 0.01 0.1 1 −2 0 2 10−3 0.01 0.1 −2 0 2 0.01 0.1 y

( k e V ) 2 0 5 1A 1742-294(D) 5 5 5 1 0 5 c c c c 0 1 h h h h 0 1 a a a a 0 1 n n n n 0 n n n n e e e e l l l l

e e e e n n n n e e e e r r r r g g g g y y y y

( ( ( ( k k k k e e e e 2 V V V V 0 1E 1743.1-2843(A2) ) ) ) ) 2 AXJ1749.1-2733 (D) 2 0 2 0 KS 1741-293(A1) 0 1A1742-294 (A2) 5 5 5 5 0 0 0 0 B.2 The XIS spectra ﬁtted with the template model of the west region

130 Figure B.2: The XIS and HXD-PIN unfolded spectra per unit solid angle.

131 Figure B.3: Same as FigureB.11

132 Figure B.4: Same as FigureB.11

133 Figure B.5: Same as FigureB.11

134 Figure B.6: Same as FigureB.11

135 B.3 The HXD-PIN spectra subtracted the contaminations from the known hard X-ray sources.

136 Figure B.7: The HXD-PIN spectra subtracted the contaminations from the known hard X-ray sources.

137 Figure B.8: Same as the Figure B.7

138 Figure B.9: Same as the Figure B.7

139 Figure B.10: Same as the Figure B.7

140 B.4 The XIS and HXD-PIN Unfolded Spectra

141 Figure B.11: The XIS and HXD-PIN unfolded spectra per unit solid angle.

142 Figure B.12: Same as FigureB.11

143 Figure B.13: Same as FigureB.11

144 Figure B.14: Same as FigureB.11

145 Figure B.15: Same as FigureB.11

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149 Acknowledgments

First of all, I would like to thank very much my supervisor Prof. T. Takahashi for his continuous guidance throughout the ﬁve years of my graduate course. I am also thankful to Dr. Motohide Kokubun for the precise comments on the Galactic center and the data analysis. I thank Dr. Shin Watanabe and Dr. Yasunobu Uchiyama for their advice and helpful comment to this theis I also appreciate all members in Takahashi group and the Suzaku HXD team. l really learned a lot of things from the powerful and helpful colleagues. Finally, I thank my family for their support and understanding.

150