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Preface
The Japanease astrophysical payload MAXI (Monitor of All-sky X-ray Image)has been monitoring the variable X-ray sky since August 15, 2009 from theInternational Space Station. MAXI is an advanced all-sky monitor in the new
century. During the seven years, some discoveries were made with MAXI.
This decade was spacialized by the all-sky monitors starting from Swift/BAT,
followed by Fermi/LAT, and then by MAXI. There are also various monitoring
programs operating on ground from optical to radio bands. The multi-wavelength
observations have been established. It is a new astrophysics which cannot be
achieved by snapshot telescope observations. Moreover, thanks to the Internet,
those instuments are connected online, which has realized the time-domain
astronomy. With these, we mankind can trace the variability in many wavelength
after emergence of a transient. Recent detections of the gravitational waves havesignificantly increased the importance of MAXI as a currently operating all-sky
monitor, and as a member of multi-messenger astronomy which covers not only
electromagnetic waves, but also neutrinos, and gravitational waves.
MAXI has detected about a hundred of gamma-ray bursts, and been performing
unbiased watch for tens of stellar flares. In addition to them, long-term X-ray
variations of about a hundred of sources are continuously monitored. MAXI has
discovered 17 new X-ray sources. By being followed up by other telescopes, theirnature were revealed. One of them was the histroical witness of the nuclear ignitionof a massive nova, six were black-hole binaries, and five with neutron star binaries.
Nevertheless, four of them remain totally unidentified, and are considered to form apotentially new class of short soft X-ray transients.
This conference was held to review the successful seven years of observations of
MAXI. It also included Hitomi results and some newmissions in China and India.
The conference was held on December 5 – 7, 2016 in RIKEN Wako campus. The
participants were 143 from 13 countries, including 28 foreign participants. We had
48 oral presentations and 61 posters. Their contents are summaried in thisproceeding.
This is the fifth in a series of the international MAXI workshops. It wasorganized by RIKEN (The Institute of Physical and Chemical Research). We
acknowledge the support by the RIKEN symposium, and the Society of Promotion ofSpace Science (Japan). Also we appreciate all the participants, who made theconference successful. Finally, Prof. Masaru Matsuoka and Prof. Hiroshi Tsunemi
will retire in March 2017. Prof. Matsuoka has worked as the MAXI PI, and Prof.
Tsunemi has worked as the MAXI/SSC PI. We deeply admire their continuous
efforts which have led MAXI to be a histrocal successful mission.
Chair of Organizers
Tatehiro Mihara

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Color versions of the proceedings are available at
http://maxi.riken.jp/conf/sevenyears/pdf/
MAXI Highlights
MAXI 7 Years Highlights

3

Tatehiro Mihara and the MAXI teamMAXI-SSC Results In 7 Years

9

Hiroshi Tsunemi, Hiroshi Tomida, Satoshi Nakahira, and SSC team

Discovery of 17 X-ray Transients with MAXI/GSC and their Nature

15

Hitoshi Negoro and MAXI teamMAXI Data Archive Plan

21

K. Ebisawa, S. Nahahira, Y. Sugawara, S. Ueno, H. Tomida, T. Furusho, T. Mihara, M. Sugizaki,
M. Shidatsu, M. Serino, H. Negoro, Y. Nakagawa, and L. Angelini

The MAXI/GSC catalog in the low Galactic-latitude sky

27

Takafumi Hori, Yoshihiro Ueda, Megumi Shidatsu, Taiki Kawamuro, Atsushi Tanimoto, Saeko Oda,
and the MAXI team

MAXI/GSC 7-year all-sky map and studies of large-scale diﬀuse structures

29

Megumi Shidatsu, Yoshihiro Ueda, Taiki Kawamuro, Takafumi Hori, Satoshi Nakahira, and the MAXI teamOperation and Calibration of MAXI GSC Xe Gas Counters on the ISS for over 7 years

31

Mutsumi Sugizaki, on behalf of MAXI GSC TeamThe Current Status of Solid-state Slit Camera aboard MAXI

33

Hiroshi Tomida, Hiroshi Tsunemi, Satoshi Nakahira, and MAXI/SSC team

Energy spectra of bright clusters of galaxies obtained from the MAXI data archive

35

Tae Furusho, Yasuharu Sugawara, Satoshi Nakahira, and Ken Ebisawa

MAXI Transients
The Swift and MAXI Galactic Transient Collaboration

39

Jamie Kennea, Hitoshi Negoro, Phil Evans, Andrew Beardmore, Hans Krimm, Motoko Serino,
Patrizia Romano, and Kazutaka Yamaoka

The orbital period of MAXI J1305−704

45

Aarran W. Shaw, Philip A. Charles, Jorge Casares, and Danny Steeghs

Bright novae— Indications of the spectrum of MAXI J0158−744 —49

Toshikazu Shigeyama, Kentaro Wada, and Yukari Ohtani

Classiﬁcation of Hard X-ray Light Curves of Black Hole Binary Transients53

K. Ebisawa, S. Nahahira, T. Amagasa, F. Hayashi, Jun Fan, J. Sugimoto, and A. Yoshida

Optical Monitoring of the MAXI sources by the optical telescope AROMA-N

55

Yoshiki Kitaoka, Takanori Sakamoto, and Atsumasa Yoshida

A panchromatic overview of accreting binary systems and their associated relativistic jets

57

Stephane Corbel

Search for soft X-ray ﬂashes at the ignition phase of optical novae using GSC data

59

Mikio Morii

Stellar Mass Black Holes
Recent Constraints on Jet physics and Properties Obtained from High Energy Observations ofMicroquasars
Jérôme Rodriguez

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63

Violent Optical Variations and Correlation between Optical and X-ray Variability During the2015 Outbursts in V404 Cygni

69

Mariko Kimura, Keisuke Isogai, Taichi Kato, Yoshihiro Ueda, Satoshi Nakahira, Megumi Shidatsu,
Teruaki Enoto, Takafumi Hori, Daisaku Nogami, Akira Imada, Colin Littleﬁeld, Ryoko Ishioka,
Elena P. Pavlenko, Alexei S. Pozanenko, Hyungsuk Tak and the Variable Star Network (VSNET) teamTwo optical emission components with diﬀerent variability in V404 Cygni

73

Y. Tachibana, T. Yoshii, H. Hanayama, and N. Kawai

Observation of the long-term variability of Cygnus X-1 with MAXI

77

Juri Sugimoto, Shunji Kitamoto, Tatehiro Mihara, and the MAXI teamAccretion Flow Properties of three MAXI Black Hole Candidates: Analysis with the TCAF
Solution81

Dipak Debnath, Sandip K. Chakrabarti, Arghajit Jana, D. Chatterjee, A. A. Molla, and S. Mondal

Inﬂow-Outﬂow Properties of Accretion Disk around MAXI J1836−194 with TCAF Solution
during its 2011 Outburst

87

Arghajit Jana, Dipak Debnath, S. K. Chakrabarti, D. Chatterjee, A. A. Molla, and S. Mondal

Enhanced power spectrum analysis of short-term X-ray variability using MAXI/GSC data

89

Tomofumi Kawase and Hitoshi Negoro

Reinvestigation of the relation between the state transition and the mass accretion rate in blackhole candidates

91

Takahiro Masumitsu and Hitoshi Negoro

A model for the ‘outlier’ radio/X-ray correlation in black hole X-ray binaries

93

Erlin Qiao

Suzaku Observation of the Eastern X-ray Lobe of SS 433

95

Hideki Uchiyama and Yuta Kanbe

The broad-band X-ray power spectrum in black hole binaries: from months to milliseconds

97

Magnus Axelsson

X-ray monitering of the Galactic microquasar V4641 Sgr98

Kazutaka Yamaoka

Low Mass X-ray Binaries and Superbursts
Spectral states in NS-LMXBs observed with MAXI/GSC and Swift/BAT

103

Kazumi Asai, Tatehiro Mihara, Masaru Matsuoka, and Mutsumi Sugizaki

Overview of Spectral Change in NS-LMXB

107

Hiromitsu Takahashi

Unveiling the spectral transition of Aql X-1 from the hard to soft state111

Ko Ono, Kazuo Makishima, Soki Sakurai, Zhongli Zhang, Kazutaka Yamaoka, and Kazuhiro Nakazawa

The MAXI/GSC X-ray Burst Catalog

115

Takanori Sakamoto, Hikaru Ohtsuki, Motoko Serino, Atsumasa Yoshida, and MAXI teamHard X-ray study of low mass X-ray binary Sco X-1 with Suzaku

117

Chiho Okada, Hiromitsu Takahashi, Tatsuya Nakaoka, Takao Kitaguchi, Yasushi Fukazawa,
Tadayasu Dotani, Yoshitomo Maeda, Satoru Katsuda, Hideyuki Mori, and Murakami Hiroshi

Black holes and Neutron stars in Low Mass X-ray binaries

119

Chris Done

Understanding superbursts

121Jean in ’t ZandMAXI Detections of Superbursts— Updates on Long X-ray Bursts —129

Motoko Serino and Wataru Iwakiri

MAXI/GSC observations of new superbursts from Ser X-1 and 4U 1705−44
Wataru B. Iwakiri, Motoko Serino, and Jean in ’t Zand

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133

Binary X-ray Pulsars
Monitoring X-ray Binary Pulsars with Fermi/GBM

137

P. A. Jenke and C. A. Wilson-Hodge

Luminosity and spin-up relation in X-ray binary pulsars with MAXI/GSC and Fermi/GBM

143

Mutsumi Sugizaki

MAXI/GSC observation of X-ray outbursts from Be/X-ray binary pulsars

149

Motoki Nakajima, on behalf of the MAXI team

Cyclotron absorption lines in the era of Suzaku and NuSTAR

153

Gaurava K. Jaisawal and Sachindra Naik

X-ray emission of the Be star/pulsar system PSR B1259−63 at non-periastron phases

157

Yuki Ono, Yoichi Yatsu, Mikio Morii, Nobuyuki Kawai, and the MAXI teamLong-term variability of Cen X-3 with MAXI

159

José Joaquı́n Rodes, Tatehiro Mihara, José Miguel Torrejón, Satoshi Nakahira, Guillermo Bernabéu,
and Graciela Sanjurjo

Orbital resolved spectroscopy of GX 301−2 with MAXI-GSC

161

Nazma Islam and Biswajit PaulSFXTs vs sgHMXBs: does the diﬀerence lie in the companion wind?

163

Pragati Pradhan, Biswajit Paul, and Enrico Bozzo

Indications for a long-term periodicity in EXO 2030+375

165

Eva Laplace, Tatehiro Mihara, Yuki Moritani, Motoki Nakajima, Toshihiro Takagi, Kazuo Makishima,
and Andrea Santangelo

MAXI Observations of Supergiant Fast X-ray Transients

167

Ai Sakamaki and Hitoshi Negoro

Monitoring Supergiant Fast X-ray Transients with Swift and XMM-Newton

169

Enrico Bozzo

Discovery of keV excess emission in the isolated neutron star RX J1856.5−3754

171

Tomokage Yoneyama

Stellar Flares
Stellar ﬂares detected with MAXI for the past 7 years

175

Yohko Tsuboi, Ryo Sasaki, Yumiko Nakamura, Ken Yabuki, Atsushi Kawagoe, Soichiro Kaneto,
Yasuharu Sugawara, Hitoshi Negoro, and the MAXI Team

Transient Magnetars and Magnetic Field Evolution of Neutron Stars

179

Teruaki Enoto

Suzaku/WAM Hard X-ray Observations of Solar Flares— from 2005 to 2014 —185

Seiya Yabe, M. S. Tashiro, Y. Terada, T. Yasuda, K. Yamaoka, K. Makishima, W. Iwakiri, M. Yamauchi,
N. Ohmori, S. Sugita, Y. Fukazawa, M. Ohno, T. Kawano, T. Takahashi, M. Kokubun, Y. Nakagawa,Y. Urata, K. Nakazawa, on behalf of the Suzaku/WAM teamThe X-ray luminosities in quiescent phase for superﬂare stars

187

Ken Yabuki, Yohko Tsuboi, Ryo Sasaki, Yasuharu Sugawara, Satoru Katsuda, and Hiroyuki Maehara

Statistical research of hyper X-ray ﬂares detected with MAXI
— a correlation between the quiescent luminosity and the largest ﬂare energy —189

Ryo Sasaki, Yohko Tsuboi, Satoru Katsuda, Ken Yabuki, Yumiko Nakamura, Yasuharu Sugawara,
Masaru Matsuoka, and the MAXI teamSuper ﬂare candidates detected by XMM-Newton

191

Y. Nakamura, Y. Tsuboi, Y. Terashima, Y. Sugawara, and S. Katsuda

X-ray spectral variability of the colliding wind binary WR140— the origin of the cool plasma component —Yasuharu Sugawara, Yoshitomo Maeda, Kenji Hamaguchi, and Yohko Tsuboix193Superﬂares on Sun-like stars

195

Kazunari Shibata

Active Galactic Nuclei
Hard X-ray Luminosity Function of Tidal Disruption Events:
Extragalactic Survey

First Results from MAXI

199

Taiki Kawamuro, Yoshihiro Ueda, Megumi Shidatsu, Takafumi Hori, Nobuyuki Kawai, Hitoshi Negoro,
and Tatehiro MiharaFirst Statistical Tests for Clumpy-Torus Models: Constraints from RXTE Monitoring of Seyfert
AGN

205

Alex Markowitz, Mirko Krumpe, and Robert Nikutta

A Novel View of AGN Accretion Flows Revealed by X-ray and Optical Monitoring

211

Kazuo Makishima, Hirofumi Noda, Takeo Minezaki, Mitsuru Kokubo, Tomoki Morokuma, Mamoru Doi,
Yuki Moritani, Kazuhiro Nakazawa, Makoto Watanabe, Hikaru Nakao, Masataka Imai, Kenji Kawaguchi,
Katsutoshi Takagi, Koji S. Kawabata, Tatsuya Nakaoka, Miho Kawabata, Makoto Uemura,
Yasushi Fukazawa, Michitoshi Yoshida, Kumiko Morihana, Yoich Itoh, Yuhei Takagi, Akira Arai,
Ryosuke Itoh, Yoshihiko Saito, and Shin’ya Yamada

An X-ray spectral variability of fast disk winds in AGN

215

Kouichi Hagino, Hirokazu Odaka, Chris Done, Shin Watanabe, and Tadayuki Takahashi

Revisiting the cosmological evolution of X-ray selected blazars with the Swift/BAT 70-month allsky survey data

221

Koyo Toda, Yoshiyuki Inoue, Yasuyuki Tanaka, and Yasushi Fukazawa

MAXI and AKARI investigation of nearby active galactic nuclei

223

Naoki Isobe, Taiki Kawamuro, Shinki Oyabu, Takao Nakagawa, Shunsuke Baba, Kenichi Yano,
Yoshihiro Ueda, and Yoshiki Toba

Kanata optical monitoring of radio galaxy NGC 1275

225

Kensei Shiki, Y. Fukazawa, Y. Tanaka, R. Itoh, K. Takaki, M. Kawabata, T. Nakaoka, T. Abe,N. Kawahara, and J. Hirochi

X-Ray Spectral Variability of the Narrow-Line Seyfert 1 Galaxy NGC 4051 Observed with Suzaku

227

Manami Seino, Hirofumi Noda, Katsuma Miyake, Kazuhiro Nakazawa, and Kazuo Makishima

The new primary X-ray spectral component of IC4329A conﬁrmed with Suzaku and NuSTAR

229

Katsuma Miyake, Hirofumi Noda, Shin’ya Yamada, Kazuo Makishima, and Kazuhiro Nakazawa

Variability Analysis of X-ray spectrum of Blazar OJ 287 from Suzaku/XIS and Swift/XRT(2005–2016)

231

Siti Fatima

Synergetic Growth of the Rayleigh-Taylor and Richtmyer-Meshkov Instabilities in the Relativistic
Jet

232

Jin Matsumoto

Multi-wavelength observation of blazar Mrk 421 in extreme X-ray ﬂare
Ryosuke Itoh

xi

233

Gravitational Wave Sources
Binary black hole remnants of ﬁrst stars for the gravitational wave source

237

Tomoya Kinugawa

Search for X-ray Counterparts of Gravitational Wave Events

241

Nobuyuki Kawai, Motoko Serino, Hitoshi Negoro, on behalf of the MAXI Team

X-ray Transient Localization
counterpart searchExperiment

aboarda

micro-satellitefor

multi-messenger

243

Tatsuya Sawano, Daisuke Yonetoku, Kazuki Yoshida, Yasuaki Kagawa, Kentaro Kawai, Masao Ina,
Tatehiro Mihara, and Hirokazu Ikeda

Explore the Transient X-ray Sky with Einstein Probe

247

Weimin Yuan

Development of the Hard X-ray Monitor onboard WF-MAXI

255

Satoshi Sugita, Shohei Harita, Yoichi Yatsu, Naoki, Isobe, Nobuyuki Kawai, Makoto Arimoto,
Hirokazu Ikeda, Hiroshi Tomida, Shiro Ueno, Tatehiro Mihara, Motoko Serino, Hiroshi Tsunemi,
Atsumasa Yoshida, Takanori Sakamoto, Takayoshi Kohmura, Hitoshi Negoro, Yoshihiro Ueda,
Yohko Tsuboi, and Ken Ebisawa

Advanced LIGO First Light: Multimessenger Astrophysics at the Birth of Gravitational-Wave
Observatory

257

Leo Singer

Gamma-ray Bursts
Gamma-ray bursts with MAXI and CALET

261

Satoshi Nakahira, on behalf of the CALET teamLocalization of Gamma-ray Bursts with BGO Active Shield of the Soft Gamma-ray Detector andHard X-ray Imager onboard Hitomi

267

Koji Tanaka, Masanori Ohno, Yasushi Fukazawa, Tsunefumi Mizuno, Hiromitsu Takahashi,
Takao Kitaguti, Yasuyuki Tanaka, Jun’ichiro Katsuta, Takafumi Kawano, Hiroyuki Shirakawa,
Ikumi Edahiro, Sho Habata, Chiho Okada, Norie Ohashi, Tadayuki Takahashi, Motohide Kokubun,
Shin Watanabe, Goro Sato, Masayuki Ohta, Koichi Hagino, Rie Sato, Kunishiro Mori,
Yusuke Uchida, Miho Katsuragawa, Ryota Tomaru, Hiroki Yoneda, Kazuhiro Nakazawa, Shunsuke Torii,
Soki Sakurai, Makoto Sasano, Hidetoshi Nishida, Shogo Kobayashi, Hiroaki Murakami, Ko Ono,
Yuichi Kato, Kastuma Miyake, Yoshihiro Furuta, Yuki Murota, Hiroyasu Tajima, Masaomi Kinoshita,
Kazutaka Yamaoka, Katsuhiro Hayashi, Jun Kataoka, Tatsuhiko Saito, Masaki Yoshino, Takehiro Mimura,
Yoichi Yatsu, Shin’ya Saito, Hideki Uchiyama, Takeshi Nakamori, Shin’ichiro Takeda, Toshio Nakano,Takayuki Yuasa, Yuto Ichinohe, Daisuke Yonetoku, and the Hitomi HXI/SGD teamSuzaku Wide-band All-sky Monitor view of soft gamma-ray transient269

Masanori Ohno, Kazutaka Yamaoka, Makoto S. Tashiro, Kevin Hurley, Hans A. Krimm, Amy Y. Lien,
Norisuke Ohmori, Satoshi Sugita, Yuji Urata, Tetsuya Yasuda, Ryuuji Kinoshita, Yusuke Nishioka,
Yoshitaka Hanabata, Jun’ichi Enomoto, Takeshi Fujinuma, Yasushi Fukazawa, Wataru Iwakiri,
Takafumi Kawano, Motohide Kokubun, Kazuo Makishima, Shunsuke Matsuoka, Tsutomu Nagayoshi,
Yujin Nakagawa, Souhei Nakaya, Kazuhiro Nakazawa, Takanori Sakamoto, Tadayuki Takahashi,
Sawako Takeda, Yukikatsu Terada, Seiya Yabe, and Makoto Yamauchi

Energy response function of CALET Gamma ray Burst Monitor

271

Y. Yamada, T. Sakamoto, A. Yoshida, for the CALET collaboration

Overview and ﬁrst year observation of CALET Gamma ray Burst Monitor

273

Yuta Kawakubo, Atsumasa Yoshida, Takanori Sakamoto, for the CALET collaboration

Extremely Soft X-ray Flash as the indicator of oﬀ-axis orphan GRB afterglow

275

Yuji Urata, Kuiyun Huang, Ryo Yamazaki, and Takanori Sakamoto

Gamma-ray burst observation for eight and half year by Akeno MITSuME telescope
Taichi Fujiwara, Yoshihiko Saito, Ryosuke Itoh, Yutaro Tachibana, Yuki Ono, Kotaro Morita,
Kouta Saisho, and Nobuyuki Kawai

xii

277

The ﬁrst Suzaku/WAM Gamma-ray Burst catalog

279

Norisuke Ohmori, Kazutaka Yamaoka, Masanori Ohno, Satoshi
Makoto Yamauchi, Kevin Hurley, Yuji Urata, and Makoto Tashiro

Sugita,HirakuYoshida,SGRB rate from Yonetoku relation for SGRB and the detectability of GW in O2 of aLIGO/aVirgo

281

Takashi Nakamura

GRB radio view from all angles

282

Alexander J. van der Horst

POLAR – The gamma-ray burst polarimetry experiment onboard China’s TG-2 spacelab

283

Shuang-Nan Zhang

Hitomi
Hitomi observation of the Perseus cluster— gas motions and resonant scattering —287

Kosuke Sato, Maki Furukawa, Toru Sasaki, Kyoko Matsushita, Frits Paerels, Irina Zhuravleva,
Anna Ogorzalek, Norbert Werner, Steve Allen, Hiroya Yamaguchi, Greg Brown, Shinya Nakashima,
Yasushi Fukazawa, Masanori Ohno, Tetsu Kitayama, Brian McNamara, and the Hitomi collaboration

Hitomi observation of the Perseus cluster— temperature and elemental abundances —288

Shinya Nakashima, Hiroya Yamaguchi, Kosuke Sato,
Marshall Bautz, and the Hitomi Collaboration

Kyoko Matsushita,

Aurora Simionescu,

Hitomi Results of NGC 1275: The Origin of Fe-Kα Line289

Hirofumi Noda, Yasushi Fukazawa, Koichi Hagino, Masanori Ohno, Masahiro Tsujimoto,Takaaki Tanaka, Hiroyuki Uchida, Taiki Kawamuro, Francesco Tombesi, Scott Porter, Caroline Kilbourne,
Laura Brenneman, Christpher Reynolds, Richard Mushotzky, and the Hitomi collaboration

Hitomi Observations of Supernova Remnants N132D and G21.5-0.9

291

Satoru Katsuda, Hiroya Yamaguchi, Makoto Sawada, Eric D. Miller, Kumiko Nobukawa,
Masayoshi Nobukawa, Takaaki Tanaka, Hiroshi Nakajima, Aya Bamba, John P. Hughes,
Masahiro Tsujimoto, Toshiki Sato, Kosuke Sato, Lorella Angelini, Benson Guest, Hiroyuki Uchida,
Hiromitsu Takahashi, Yukikatsu Terada, Shu Koyama, Kenya Oshimizu, Caroline A. Kilbourne,
Katsuji Koyama, Samar Saﬁ-Harb, Brian J. Williams, and the Hitomi Collaboration

Atomic data and spectral modeling constraints from high-resolution X-ray spectroscopic
observations of the Perseus cluster with Hitomi

292

M. Sawada, J. Kaastra, L. Gu, H. Akamatsu, J. de Plaa, C. de Vries, R. Smith, A. Foster, H. Odaka,
M. Leutenegger, M. Markevitch, C. Kilbourne, H. Yamaguchi, G. Brown, F. Paerels, N. Ota, T. Tamura,
S. Nakashima, S. Ueda, K. Matsushita, K. Sato, R. Mushotzky, T. Ohashi, and Hitomi Collaboration

Hitomi Observation of the Highly Obscured High-Mass X-ray Binary IGR J16318−4848

293

Hiroshi Nakajima, Kiyoshi Hayashida, Tim Kallman, Caroline Kilbourne, and the Hitomi Collaboration

The Hard X-ray Imager onboard Hitomi (ASTRO-H)

295

Kazuhiro Nakazawa, Goro Sato, Motohide Kokubun, Teruaki Enoto, Yasushi Fukazawa, Kouichi Hagino,
Atsushi Harayama, Katsuhiro Hayashi, Jun Kataoka, Junichiro Katsuta, Philippe Laurent,
Francois Lebrun, Olivier Limousin, Kazuo Makishima, Tsunefumi Mizuno, Kunishiro Mori,
Takeshi Nakamori, Toshio Nakano, Hirofumi Noda, Hirokazu Odaka, Masanori Ohno, Masayuki Ohta,
Shinya Saito, Rie Sato, Hiroyasu Tajima, Hiromitsu Takahashi, Tadayuki Takahashi, Shin’ichiro Takeda,
Yukikatsu Terada, Hideki Uchiyama, Yasunobu Uchiyama, Shin Watanabe, Kazutaka Yamaoka,
Yoichi Yatsu, Takayuki Yuasa, and the HXI teamIn-ﬂight performance of the Soft Gamma-ray Detector (SGD) onboard Hitomi
Yasushi Fukazawa, Hiroyasu Tajima, Shin Watanabe, Roger Blandford, Teruaki Enoto, Andrea Goldwurm,
Kouichi Hagino, Katsuhiro Hayashi, Yuto Ichinohe, Jun Kataoka, Junichiro Katsuta, Takao Kitaguchi,
Motohide Kokubun, Philippe Laurent, François Lebrun, Olivier Limousin, Grzegorz M. Madejski,
Kazuo Makishima, Tsunefumi Mizuno, Kunishiro Mori, Takeshi Nakamori, Toshio Nakano,Kazuhiro Nakazawa, Hirofumi Noda, Hirokazu Odaka, Masanori Ohno, Masayuki Ohta, Shinya Saito,
Goro Sato, Rie Sato, Shin’ichiro Takeda, Hiromitsu Takahashi, Tadayuki Takahashi, Takaaki Tanaka,
Yasuyuki Tanaka, Yukikatsu Terada, Hideki Uchiyama, Yasunobu Uchiyama, Kazutaka Yamaoka,
Yoichi Yatsu, Daisuke Yonetoku, and Takayuki Yuasa

xiii

297

Development and in-orbit performance of all-sky monitoring function of BGO active shield of thesoft gamma-ray detector onboard Hitomi

299

Masanori Ohno, Yasushi Fukazawa, Tsunefumi Mizuno, Hiromitsu Takahashi, Takao Kitaguti,
Yasuyuki Tanaka, Jun’ichiro Katsuta, Takafumi Kawano, Hiroyuki Shirakawa, Ikumi Edahiro, Sho Habata,
Chiho Okada, Norie Ohashi, Koji Tanaka, Tadayuki Takahashi, Motohide Kokubun, Shin Watanabe,
Goro Sato, Masayuki Ohta, Koichi Hagino, Rie Sato, Kunishiro Mori, Yusuke Uchida, Miho Katsuragawa,
Ryota Tomaru, Hiroki Yoneda, Kazuhiro Nakazawa, Shunsuke Torii, Soki Sakurai, Makoto Sasano,
Hidetoshi Nishida, Shogo Kobayashi, Hiroaki Murakami, Ko Ono, Yuichi Kato, Kastuma Miyake,
Yoshihiro Furuta, Yuki Murota, Hiroyasu Tajima, Masaomi Kinoshita, Kazutaka Yamaoka,
Katsuhiro Hayashi, Jun Kataoka, Tatsuhiko Saito, Masaki Yoshino, Takehiro Mimura, Yoichi Yatsu,
Shin’ya Saito, Hideki Uchiyama, Takeshi Nakamori, Shin’ichiro Takeda, Toshio Nakano, Takayuki Yuasa,
Yuto Ichinohe, Daisuke Yonetoku, and the Hitomi HXI/SGD teamInstrumentation
X-ray Polarimetry Mission PRAXyS

305

Toru Tamagawa, on behalf of the PRAXyS teamEnd-to-End Test of the Mirror and Polarimeter of PRAXyS

307

Toshio Nakano, Toru Tamagawa, Wataru Iwakiri, Asami Hayato, Megu Kubota, Kazuki Nishida,
Takao Kitaguchi, Koyo Toda, Noriyuki Tanaka, Tsunefumi Mizuno, Teruaki Enoto, Ikuyuki Mitsuishi,
Yuzuru Tawara, Shota Inoue, Kiyoshi Hayashida, Keith Jahoda, Joe Hill, Ross McCurdy,
and PRAXyS teamiWF-MAXI : Soft X-ray Transient Monitor on the ISS

309

N. Kawai, Y. Yatsu, N. Isobe, S. Sugita, S. Harita, H. Tomida, S. Ueno, K. Ebisawa, T. Dotani, T. Mihara,
M. Serino, M. Arimoto, H. Tsunemi, T. Sakamoto, T. Kohmura, H. Negoro, Y. Ueda, T. Nakamura,T. Tanaka, Y. Tsuboi, M. Morii, M. Yoshida, N. Kanda, and M. Vagins

A broadband X-ray imaging spectroscopy with high-angular resolution: the FORCE mission

311

Koji Mori, Takeshi Go Tsuru, Kazuhiro Nakazawa, Yoshihiro Ueda, Takashi Okajima, Hiroshi Murakami,
Hisamitsu Awaki, Hironori Matsumoto, Yasushi Fukazawa, Hiroshi Tsunemi, and Tadayuki Takahashi

Development of X-ray SOI pixel sensors for future X-ray wide ﬁeld cameras

313

Hideki Hayashi, T. G. Tsuru, T. Tanaka, H. Uchida, H. Matsumura, M. Itou, S. Ohmura, Y. Arai,
I. Kurachi, K. Mori, A. Takeda, Y. Nishioka, N. Takebayashi, S. Yokoyama, T. Kohmura, K. Tamasawa,
K. Negishi, K. Oono, K. Yarita, S. Nakashima, S. Kawahito, K. Kagawa, K. Yasutomi, H. Kamehara,
and S. Shrestha

Micro-satellite for Astrometry and Photometry of Gravitational Wave Sources

315

S. Harita, T. Ozawa, T. Yoshii, Y. Yatsu, N. Kawai, T. Tawara, M. Matsunaga, N. Tominaga, S. Kulkarni,
P. Bilgi, T. Sakamoto, N. Vasquez

Development of the Solar Neutron Monitor on the ChubuSat-2 Satellite

317

Kazutaka Yamaoka, Takuya Miyazawa, Shinji Oseki, Yasunobu Babazaki, Yuki Hayashi, Daisuke Ishihara,Kazuya Itoh, Hidehiro Kaneda, Hiroaki Kawahara, Shin Kubo, Koji Matsushita, Kikuko Miyata,
Hosei Nagano, Yasutaka Narusawa, Masaki Nishino, Sosuke Noda, Daeil Park, Masaaki Sadamoto,
Hiroyasu Tajima, Keisuke Tamura, Hidetaka Tanaka, Dao Ngoc Hanh Tam, Toyoki Watabe,
and Qidong Zhou

Development of an in-orbit radiation environment monitor CUBES onboard a small satellite
MIST

319

Nagomi Uchida, Tatsuya Nakaoka, Hiromitsu Takahashi, Tsunefumi Mizuno, Yasushi Fukazawa,
Mark Pearce, and Theodor-Adrian Stana

Some early results from Astrosat323

Biswajit PaulThe Hard X-ray Modulation Telescope Mission

324

Shuang-Nan Zhang

xiv

xv

1

MAXI Highlights

2

3

MAXI 7 Years Highlights
Tatehiro Mihara,1 and the MAXI team
1 MAXI team, RIKEN, Wako, Saitama, Japan

E-mail (TM): tmihara@riken.jp
ABSTRACT

Since MAXI started observation on August 15, 2009, MAXI is monitoring the X-ray sky for morethan 7 years. The light curves of about 350 sources are updated in every 4 hours or 1 day and open
through MAXI homepage (http://maxi.riken.jp). Alerts from MAXI on new X-ray transients are sent
to subscribers in the MAXI alert mailing lists at 12 seconds after the emerge as the earliest case.
Many of them were followed up by the Swift/XRT and the nature of the source were revealed.
Seventeen new sources were found by MAXI. Six of them were new black holes (BH). Soft X-ray flash
at an ignition of a nova was detected for the first time in 2011 November 11. Tidal disruption events
have formed a new category of transients after the discovery of Swift J1644+57. MAXI was usefulthat it had motnitored the source activitity before the event.
MAXI catalogs of both GSC and SSC were issued listing new members of X-ray emitters in the 21st
century. Fluxes of X-ray binaries are monitored continuously recording histories of activities of theneutron stars (NS) and BH. Monitoring data is used for detailed observations with Suzaku, Chandra
and XMM-Newton. Data with uniform quality are used for power-spectrum-density analysis in AGNs
and study of accretion disk instabilities in NS-LMXB and BH binaries. Normal and giant outbursts of
Be X-ray binaries are monitored enabling a detailed discussion on the Be stellar disk together with
optical emission-line profiles. MAXI detected large flares from blazars as well as those from active
stars, and one-hour lasting superbursts from NS-LMXB. Gamma-ray bursts have been detected
steadily, whose number reached 85. Still, we had 7 MAXI unidentified short soft transients (MUSST).
In 2015 LIGO started operation and found the gravitational wave sources. NICER, a new X-ray
telescope on the ISS, will be mounted in early 2017. Together with these new instruments, MAXI will
open new windows to understand the universe further.
KEY WORDS: MAXI, All-sky monitor, Transients, Time-domain astronomy
1.

MAXI instruments

On 2009 July 16, MAXI was launched by the
space shuttle Endevour (Matsuoka et al. 2009). A
Japanese astronaut Koichi Wakata moved therobotic arm to mount MAXI to JEM-EF port #1(Figure 1). MAXI/GSC (Mihara et al. 2011, Sugizaki et al. 2011) has two field of views (FOV),
Horizontal and Zenithal, each with 160o3o FOV.
Since the operation of GSC is limitted to the low
background region around the equator, the fraction of the working time is ~0.4 (Sugizaki et al.
2017). Even though, MAXI/GSC can scan almost
the whole (90%) sky in one orbit (92 minutes)
thanks to the two FOVs. Dwell time of one MAXI
scan is ~60 s, and the 5sensitivities are ~130mCrab / 1 scan, ~20 mCrab / 1 day, ~4 mCrab / 1
month, and ~1 mCrab / 1 year (Negoro et al.
2016).

Figure 1. Real MAXI on ISS. Photo taken from thedocked space shuttle (©NASA). Right-hand side is
the moving direction of ISS. MAXI has two FOVs.

4

Figure 2. First light of MAXI/GSC on 2009 August
15. This image was press-released by JAXA and
Riken. Two scan poles with 10o radius move on=23o lines with a presession period of the ISS orbit
(70 days).

2.

Press releases, papers and awards

There are 13 press-releases from MAXI starting with the first light of MAXI/GSC (Figure 2).
The first X-ray nova for MAXI, XTE J1752-223,
was press-released from the Astronomical Societyof Japan in 2010 September as “slow-eating BH”(Nakahira et al. 2010). Press-releases of scientificresults followed as a BH nova MAXI J1659-152(Yamaoka et al. 2012), a slow X-ray pulsar MAXI
J1409-619 (Yamaoka et al. 2011), a tidal disruption event Swift J1644+57 (Burrows et al. 2011),
Cygnus super bubble (Kimura et al. 2013), a
massive nova ignition MAXI J0158-744 (Morii etal. 2013), a nearby monster GRB 130427A (Maseli
et al. 2014), and a optical quick (~2 minutes)
followup of V404 Cyg (Kimura et al. 2016).
Refereed papers using MAXI data reached 192(as of 2017 February). Those by MAXI team are
62. Praising ~100 citations, the PASJ Excellent
Paper Award 2013 was given to Matsuoka et al.
(2009) by Astronomical Society of Japan (Figure3). We had three MAXI conferences in 2008, 2010,
and 2016. Admiring the MAXI leading of thecurrent time-domain astronomy and multimessenger astronomy, ISS R&D Conference 2016
Innovation Award in Earth and Space Science
was given to the MAXI by CASIS (Center for the
Advancement of Science In Space), American
Astronautical Society, and NASA.

Figure 3. PASJ Excellent Paper Award 2013 to the
MAXI misison paper, Matsuoka et al. (2009).

3.

Scientific results

3.1. MAXI all-sky map and catalog
The GSC all-sky map was presented in Mihara(2014). The SSC all-sky map was published by
Kimura et al. (2013). The 37 months MAXI/GSC
catalog in high galactic region was by Hiroi et al.
(2013). In there, 500 sources were detected with ahigher significance than 7. The detection limit
was ~0.4 mCrab. The log N - log S relation isconsis- tent with HEAO-1’s result in fluxes >10-11
erg cm-2 s-1. However, the 40% of object members
had changed since 1970’s. Thus MAXI has
created a catalog in the early 21st century.
3.2. MAXI transients
MAXI discovered 17 new sources in addition to
85 gamma-ray bursts (Negoro et al. 2017). The
breakdown is 1 white dwarf binary, 6 NS, 6 BH,
and 1 unknown. The completely new kind of
phenomenon was massive nova ignition MAXI
J0158-744 (Morii et al. 2013). On 2011 November
11, a bright (~ 1 Crab) but soft ( < 5 keV) X-ray
transient appeared near SMC. Fortunately both
H and Z-SSC detected it. It was not detected in
the next GSC scan (Figure 4). About 11 hours
later Swift/XRT detected a vary faint (10-3 of
MAXI’s flux) source. It was identified to a Be star(B1-2 IIIe) at SMC distance ( = 60 kpc). Then the5flux was 1040 erg s-1, which exceeds the Eddington limit for 1 Mo by a factor of 100. The GSC
spectrum can be fitted by any of power law, black
body (BB), and thermal brems. If BB was applied,
the temperature was 450 eV and the radius was
103 km. The Swift spectra can be fitted by BB
with 60
110 eV and area of 104
102 km.
From the SSC-Z observation, an ionized Ne line
was detected. Theoretical explanation followed(Ohtani et al. 2014). The size and temperature
are typical to a white dwarf (WD) and the light
curve was similar to that of a classical nova. The
difference was the faster time-scale and thehigher temperarure. The MAXI part was interpreted as an ignition phase of a nova, which wasobserved for the first time in the history of nova.
Also, it was soft X-ray against the expectation of
UV emission. The Swift part corresponds to thesuper soft source (SSS) phase. The fast time-scaleand high temperature can be explained by a very
massive (nearly Chandra- sekhar limit of) WD.
Thus the system would be a vary rare one
composed of a massive O-Ne WD and a Be star.

3.4. Monitoring X-ray binaries (XRB)
Activities of NS- or BH- XRB are monitored
once in 92 minutes and made public at MAXI HP 1
every 4 hours or every day. Periodical and nonperiodical outbursts from BeXRB were presented
by e.g. Mihara et al. (2014) and Nakajima et al.
(2017). The MAXI data are compiled in BeXRB
activity web page2 by ESA/ESAC.
The non-periodic outbursts of NS low-mass
X-ray binaries (LMXB) are studied by Asai et al.
(2012), where repetition of outbursts in several
hundred days were clearly shown. Sudden decrease in the last part of an outburst was related
to the propeller effect, from which magnetic fieldsof the NS was estimated (Asai et al. 2013). A
small outburst without spectral state transition
was named “mini outburst” in Asai et al. (2015).
It was an observational evidence of purr-type
disk instability predicted theoretically. Outburstsof BH are also watched by seveal researchers(Figure 5).
A long-term monitoring of an XRB is also useful for detailed observations. Chandra observation by Heinz et al. (2015) showed giant X-ray
light echoes around Cir X-1. This “Lord of theRings” was understood using the light curves of
MAXI in hundreds of days.

Figure 4. X-ray light curve of MAXI J0158-744.
After trigger by GSC-H, SSC-H and SSC-Z detectedsuccessively (Morii et al. 2013).

3.3. Huge stellar flares
MAXI detected huge stellar flares unbiasly.
Those are from nearby (~100pc) active stars: dMestars and RS CVn (and Algol) binaries. The peakluminosities are 1034-1036 erg s-1 and 1036-1039 erg
s-1, respectively. These are 103-107 times larger
than the largest solar flares, and even 1-104
times larger than the super flares of G-type stars(Tsuboi et al. 2016, Tsuboi et al. 2017).

Figure 5. X-ray light curves of BH binaries. Taken
from MAXI HP 3 . Hardness curves, intensityhardness plots and spectra are also available.

http://maxi.riken.jp/
http://integral.esac.esa.int/bexrbmonitor/
3 http://maxi.riken.jp/mxondem/outburst/BH/
12

6The pulsation was mointored with MAXI fromsome X-ray binary pulsars. The 7.67s pulsation of
4U 1626-67 is detected in each 60-day. Since
7.67s is short enough compared to the ~60 s scantime, we can get rid of uncertainty by variablebackground and source-flux in the folding process.
For the intensity of 30 mCrab, we need 60-day
integration to detect a pulsation. In the 7 years of
MAXI observation, 4U 1626-67 increased theluminosity by 20%. The spin up rate is related toa flux. In the historical view, 4U 1626-67 repeats
spin-up and down. Together with historical results, the P-dot and flux relation is well explainedby the Ghosh and Lamb (1979) theory (Figure 6).
Their n( ) function with the fastness parametercan handle both spin-up and down. The surface magnetic filed is known from the cyclotronfeature. Then we can determine the mass andradius of the NS if we know the distanceaccurately enough (Takagi et al. 2016).
Superburst is a long X-ray burst in NS-LMXB
lasting ~3000s. Normal X-ray burst is a H and He
burning, while the superburst contains a Cburning in addition to H and He. Since theintensity and duration match MAXI scanning,

Figure 6. Flux and pulse period derivative of 4U
1626-67 with MAXI and previous results. The
best-fit model is by Ghosh and Lamb (1979). In thisplot we added a flux uncertainty of 5.5% so as the2 to be 1.0. If we assume the dis- tance D=10 kpc,
we can obtain the mass M = 1.740.05 Mo, andradius R = 13.50.1 km of the NS very accurately.
The dashed and the dotted lines correspond tochanges of the radius by 1 km and the mass by 0.3
Mo, respectively. (Takagi et al. 2016).

MAXI detected 12 superbursts including intermediate bursts, among 28 ever detected (Serino
et al. 2016).
Taking the advantage of MAXI’s continuous
sampling, a long term power spectrum density(PSD) was studied. In Cyg X-1, long-term PSD
down to 10-7 Hz was obtained by Sugimoto et al(2016) for both soft state and hard state.
Two-layered accretion disk was discussed basedon it. They needed to correct observational gaps
and aliases arising from the sparse-sampling by
snap-shot observations. Premitive PSDs of manysources without corrections are publicly available
from MAXI HP4.
3.5. AGN variability
Some AGN show rapid variability. The mostvariable one is the BL Lac objects. Figure 7
shows 3-day flares of Mrk 421. X-ray fluxchanged by more than one-order-of-magnitude inseveral days. Long-term PSD of Mrk 421 has a
break at 10-5 Hz. In lower frequency than that,
the PSD extended in a power-law down to 10-8 Hz(Isobe et al. 2014).

Figure 7. top : MAXI light curve of Mrk 421.
bottom : long-term PSD of Mrk 421 (Isobe et al.
2014)4

http://maxi.riken.jp/PSD

7

3.6. Gravitational wave & unknown transients
Last year the histrical detection of the first gravitational wave (GW) was reported. GW 150914
occurred at 2015 September 14 9:50:45 UT.
Unbiased survey of MAXI is suitable for watchingan X-ray counter part of a GW. Unfortunately
MAXI was not operating at the time of the GW
detection because it was in a high backgroundregion. In 4 minutes later, MAXI resumed theoperation and 3  upper limit (2-20 keV) was
obtained as 0.1 c s-1 cm-2 ~ 30 mCrab ~ 10-9 erg s-1
cm-2 (Kawai et al. 2017). MAXI also put upper
limits in flux just before, long before and longafter the event. For the second one GW 151226
MAXI was operating at the occurrence, but justan upper limit was obtained (Serino et al. 2017).
In the seven years of unbiased operation, MAXI
has accumulated unknown soft short transients.
In spite of the Swift/XRT followups, no counterpart was detected for the seven sour- ces (Table 1).
We named it MAXI unidentified short softtransient (MUSST). In Figure 8 MAXI J1501-026
was shown as an example. It is soft : detected
only with MAXI in the soft X-ray band (2-10 keV)and not detected by Swift/BAT nor GRB detectors.
It is a short transient not detected in the next
scan of MAXI, and unidentified because ofundetection by the Swift/XRT follow-up.

Table 1. Seven MUSST sources
Name
Flux
MAXI / GRB
Date(Crab) Atel/GCN
M J1501-026
2015-8-26
0.44
A #7954
G 150428C
2015-4-28
0.16
G#17772
M J1540-158
2015-3-11
0.1
G#17568
G 140814A
2014-8-14
0.23
G#16686
M J0545+043 2014-4-12
0.2
A #6066
G 130407A
2013-4-07
0.17
G#14359
M J1631-639
2011-4-29
0.12
A #3316
We reported some sources with GRB names. But
they would not be GRB because they were soft anddid not accompany X-ray afterglows.

c/s/cm2

Figure 8. MAXI J1501-026 as an example of
MUSST. left : MAXI image at detection. It is very
bright. middle: light curves of MAXI scan in threeenergy bands. There was no flux in the hard band
(10-20 keV). right : X-ray image of Swift/XRT
followup. No source was detected in the MAXI error
region.

4.

MAXI future

MAXI’s operation was adopted untill 2018
March. In terms of the Japanese contribution toISS, JAXA continues to join the ISS program till
2024. They call it “op-3”. We MAXI team will
apply for a further 2-3 years of extension of MAXI
operation. New science points in this extension
will be a search for an X-ray counter part of GW
and revealing the MUSST. In general words,
search for the discovery space in the time-domain
astronomy (Figure 9). Here cooperation with new
instruments is a key. The JEM mission CALET isoperating since 2015 (Nakahira et al. 2017).
CALET/GBM observes GRB simultane- ouslywith MAXI. Advanced LIGO is operating since2014 to detect a GW. Currently O2 is going on.
But the most advance will be coorporation with
NICER.
NICER is a new ISS instrument which will belaunched in 2017 April. It will be attached to
ELC3 truss in the NASA part. NICER is an arrayof pointing X-ray telescopes without imaging.
Besides the main purpose of tracing the pulse
light curve of a single NS, it will be capable of aquick follow-up of a new MAXI transient. We
MAXI team and NICER team proposed the
OHMAN (On-orbit Hookup of MAXI And NICER)
project to connect the two on orbit for the
real-time follow-ups. The point is “Watch the
X-ray transient in X-ray while it is still bright”.

8

observations of ground observatories andsatellites in orbit. New phenomena (as an
ignition of a nova) and six black holes werediscovered. MAXI has opened a new era oftime-domain astronomy and of multi-messenger
astronomy with the highly-sensitive X-ray all-sky
monitor and the real-time alert. Together withnew instruments such as gravitational wave
detectors and the X-ray detector NICER on ISS,MAXI will be on the cutting edge of the X-ray
astronomy.
References
Figure 9. Rapidness of follow-up in X-ray andsensitivity. Discovery space in the time-domain
astronomy is shown. Some kinds of quick decaying
objects can be detected with the on-orbit MAXI –NICER connection. Here we assume
MAXI: When the explosion started in MAXI FOV,NICER: MAXI alert on-orbit and via-ground,
Swift: MAXI alert through GCN. Exposure of 1ks.

Besides the cooperation with new instruments,
continuing MAXI operation can make follows.
Coordinated observation of transients in
multi-wavelength in Japan and internationally.
New black hole binaries : 1-2/ yr
Huge stellar flares
Low-mass X-ray binary and super bursts
Gamma-ray bursts
Tidal disruption events
X-ray binary transients
3rd catalog in low galactic region (|b|<10)
Hori et al. (2017 and more sources.
More X-ray photon statistics. Sensitivity
reaches confusion limit (~0.5 mCrab).
The
number of sources reaches~1000.
With Light curves and variability index.
i.e. information in time-dimension, too.
Permanent MAXI archive at ISAS (Ebisawa
2017)
5.

Summary

Continuous monitoring of MAXI on ISS has
been providing basic information on variability of
X-ray sources, which is distributed freely to theworld. A real-time alert triggers many follow-up

Asai, K. et al. 2015, PASJ, 67, 92
Asai, K. et al. 2013, ApJ, 773, 117
Asai, K. et al. 2012, PASJ, 64, 128
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Ebisawa, K. et al. 2017, in this proceeding
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Heinz, S. et al. 2015, ApJ, 806, 265
Hiroi, K. et al. 2013, ApJS, 207, 36
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Kawamuro, T. et al, 2016, PASJ, 68, 58
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9

MAXI-SSC Results In 7 Years
Hiroshi Tsunemi,1 Hiroshi Tomida,2 Satoshi Nakahira,3 and SSC team1

Department of Earth and Space Science, Osaka University,
1-1 Machikaneyama, Toyonaka, Osaka 560-0043 Japan
2
Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency,
3-1-1 Yoshinodai, Chuo, Sagamihara, kanagawa, 252-5210 Japan
3
Human Spaceﬂight Technology Directorate, Japan Aerospace Exploration Agency,
2-1-1 Sengen, Tsukuba, Ibaraki, 305-805 Japan
E-mail(HT): tsunemi@ess.sci.osaka-u.ac.jp
Abstract
The MAXI SSC is a slit-camera using X-ray CCDs. Its ﬁeld of view is 1◦ × 90◦ and covers the wholesky. The CCD performance gradually degrades due to the charged particles while it has been working
since its launch in 2009. Due to the small eﬀective area of the slit, the SSC detection limit is higher thanthat of the MAXI GSC whereas its eﬀective energy range is between 0.6 keV and 7 keV. We have observed
both point sources and diﬀuse emission. We will brieﬂy review the SSC performance and its results in this7 years.
Key words: instruments: X-ray CCDs — surveys — X-rays:general

1.

SSC performance in orbit

The SSC is an X-ray slit-camera on board the MAXI
(Matsuoka et al., 2009) launched in 2009 (Tsunemi etal., 2010, Tomida et al., 2011). It consists of two identical cameras, SSC-H and SSC-Z, looking at a horizontal
direction and a zenith direction, respectively. Each camera contains 16 CCDs just below the mechanical slit that
limits the ﬁeld of view (FOV) of the SSC to be 1◦ × 90◦ .
CCD functions in parallel sum mode as one dimensional
imager. Each CCD has its own single-stage peltier cooler
with which it is cooled below -60◦ C when the camera
body is around -15◦ C. The peltier cooler achieves morethan 40◦ C temperature diﬀerence by using 1 A electric
current. The SSC camera body is connected to two radiators through a loop heat pipe. The surface of theradiators is painted in white so as to avoid to be too
shinny.
Figure 1 shows the thermal history of the CCD for
7 years. Almost all the CCDs are kept around -60◦ C
with almost no degradation in cooling system. We had
a CPU trouble in 2013 where SSC-Z was out of control
for almost one year. During this period, CCDs in SSC-H
were kept in cold condition while those in SSC-Z wereleft with no peltier cooling.
Each SSC camera receives Cu-K ﬂuorescence lines ineach CCD. Cu-K lines are generated by charged particles onto the slats collimator made of copper. We canFig. 1. Thermal history of the CCDs on the SSC in 7 years. SSC-Z
was out of order in 2013 when the CCD temperature is that of
the camera body temperature.

monitor the performance of the CCD by measuring thestandard deviation of the line proﬁle. Figure 2 shows theevolution of the SSC performance for 7 years. We seethat the SSC gradually degrades due to the charged particles. The degradation depends on the geometrical condition inside the SSC. The far-end of the CCD received
a least damage while that just below the slit received thebiggest damage. We should note that the SSC-Z in 2013
showed almost no degradation when its CCDs were left
unbiased at around -15◦ C.

10

the sources in SSC in soft band and those in ROSAT(0.1-2.4 keV). We see that there is a very clear correlation between them. Some sources, Cygnus X-1, Her X-1
etc. showing strong time variability deviate from a simple correlation. We also ﬁnd that the detection limit of
the SSC is a few mCrab.

Fig. 2. Evolution of the CCDs on the SSC-H (top ﬁgure) and on the
SSC-Z (bottom ﬁgure). The Cu-K line proﬁle (standard deviation)is shown as a function of time.

The CCDs of the SSC are designed to be sensitive to
X-rays while they are insensitive to visible photons by Alcoat on the imaging area. However, we ﬁnd that they aresensitive to infra-red, particularly at device edges. The
data obtained at daytime are heavily saturated depending on the sun direction. Therefore, we usually collectdata obtained at night time. This reduces the observation eﬃciency of the SSC by a factor of two.
2. Comparison with ROSAT
2.1. Point source observation
The SSC can cover the energy range between 0.6 and7 keV. The low energy limit is determined by the thermal noise while the high energy limit is practically determined by the Cu-K lines. Tomida et al. (2016) generated the point source catalog using the SSC data between
2010 August to 2014 April and listed 170 sources. They
measured the source intensity below and above 1.85 keV(soft band and hard band, respectively). The soft bandis out of the GSC energy range (Mihara et al., 2011).
There are 32 sources only seen in the soft band whileother 30 sources are seen only in the hard band. The
soft band data are the unique measurement of the entire
sky since the ROSAT all sky survey in 1991 (Voges et al.,1999). Figure 3 shows the intensity correlation between

Fig. 3. Intensity correlation between point sources detected by the
SSC and those by ROSAT. Some sources show intensity change
in this 20 years.

2.2. Diﬀuse structure in the entire sky
The SSC covers the entire sky in unbiased scan observation. Since its energy range is below 2 keV, this is also aunique sky map since the ROSAT all sky survey in 1991
(Freyberg et al., 1999). Figure 4 shows the comparison of the sky map between the SSC (0.6-1.0 keV) and
the ROSAT (0.44-1.21 keV). In both maps, bright point
sources are eliminated so that the we can see a diﬀuse
emission structure in detail. The angular resolution of
the SSC is about a few degrees while that of ROSAT isabout an order of magnitude better. We see that they
are quite similar to each other with taking into accountthe diﬀerence in angular resolution.
3. GRB Nova and CSB
3.1. GRBThe SSC scans the sky in the fan beam according to theorbital motion of the International Space Station orbital
motion. Therefore, a typical on-source time is about
50 seconds at every 90 minutes. We have detected several
short-time transients under this circumstance. Amongthem, we detected the initial phase of GRB100418A fromits very beginning (Imatani et al., 2016). Some GRBs
having a pre-cursor were detected by the BAT. The following observation by the Swift/XRT could cover themain part of it (Romano, et al., 2006). While many ofthem show no pre-cursor whose main peak is diﬃcult tobe observed from its beginning.

11

Fig. 5. Entire light curve of GRB100418B. The ﬁrst 50 second wasobserved by the SSC and followed by the Swift/XRT.

Fig. 4. Comparison of the SSC diﬀuse map with that of ROSAT.
With taking into account the diﬀerence in spatial resolution, we
see that they are identical to each other.

Figure 5 shows the light curve of GRB100418A. It wasa typical long GRB. When the source entered the FOVof the SSC, the GRB started that was detected by the
BAT on Swift. It reached a maximum ﬂux of a few
ten times brighter than that of the Crab nebula. The
SSC monitored its light curve for the initial 50 seconds.
The Swift/XRT pointed to the source 71 second after its
emergence (Marshall, et al., 2011).
The light curve of the GRB can be well expressed bya power law of time. By adding the SSC data in theprompt emission, we found that the light curve could be
expressed not by a power law but by a combination of
a power law and an exponential component of a decay
time about 30 second. In the ﬁrst 300 second, the exponential components dominates the emission. We foundthat GRB100418A is a typical long GRB having a very
low value of Ep(≤ 8.3 keV) that is a very rare GRB. It
is also consistent with the Yonetoku-relation (Yonetoku
et al., 2010) while it is marginal with the Amati-relation(Amati et al., 2008).
3.2. NovaMAXI detected many transients including nova. Since
the GSC has much better detection eﬃciency, the SSC
can conﬁrm it if it is bright enough. A MAXI nova,
MAXI J0158−744, was discovered in the SMC by the
GSC in 2011. The SSC followed the observation about

20 minutes after. It was still bright and we obtained
the X-ray spectrum where we detected the emission lines
from Ne and Mg (Morii et al., 2013) showing the thinthermal origin.
We conﬁrmed that the transient is a nova explosion, on
a white dwarf in a binary with a Be star. An early turnon of the super-soft X-ray source (SSS) phase, the short
SSS phase duration of about one month, and a 0.92 keV
Ne emission line found in the third MAXI scan, suggestthat the explosion involves a small amount of ejecta andis produced on an unusually massive O-Ne white dwarf
close to, or possibly over, the Chandrasekhar limit. We
propose that the huge luminosity detected with MAXI
is due to the ﬁreball phase, a direct manifestation of theignition of the thermonuclear runaway process in a nova
explosion.
3.3. Cygnus Super Bubble
There are several large scale structures in soft X-ray bandas shown in ﬁgure 4. One of them is the Cygnus Super
Bubble (CSB) that is well-known discovered by HEAO-1
in 1980 (Cash, et al., 1980) and ROSAT also conﬁrmed it.
These observations were done by using a gas proportionalcounter which was diﬃcult to reveal its spectral origin.
Figure 6 shows the SSC map around the CSB where there
are 3 Cygnus point sources and the Cygnus Loop. The
CSB is an extended source of about 15◦ × 25◦ .
We detected emission lines of Ne and Mg as shownin ﬁgure 7 which conﬁrms a thin thermal origin rather
than a non-thermal origin. The temperature is about
0.23±0.01 keV with a depleted abundance of 0.26±0.1
solar. The abundance depletion is also known in the
Cygnus Loop (Miyata et al., 1994) that is close to the
CSB. It may reﬂect the interstellar matter in this region. By adding the ROSAT data at low energy, we also
conﬁrm that the interstellar absorption feature showsno variation within the uncertainties suggesting that the12Fig. 6. CSB image in the Cygnus region. There are three points
sources (Cyg X-1, X-2 and X-3) and the Cygnus Loop.

CSB is a single unity. The Cygnus OB2 association isdiﬃcult to be an energy source, since it is oﬀ-centered of
the CSB. With taking into account the total thermal energy, we conclude that it can originate from a hypernova
event rather than a rapid sequence of many supernovae(Kimura et al., 2013).

Fig. 8. SSC band map images of the whole sky

Fig. 7. X-ray spectrum of the CSB. It can be well expressed by a thinthermal origin with a depleted abundance.

4. Diﬀuse structure and Fermi bubble
Figure 8 shows the SSC sky maps in three energy bands:
0.6-1 keV, 1-2 keV and 2-4 keV. The 2-4 keV band map
shows many point sources as well as the galactic ridge
emission that is clearly seen in the GSC map. The 0.61 keV map shows diﬀuse emission in large structure.
Figure 9 shows the SSC color map where red, green
and blue correspond to those of three bands in ﬁgure 8.
We see that low energy diﬀuse structures are seen in red,
while point sources along the galactic plane are seen inblue. The most prominent large structure is seen in thegalactic center. We see that there is a large structurein the northern hemisphere, emerging from the galactic
plane along the North Polar Spur and reaches near thenorth pole and goes down to the galactic plane. We
also see the structure in the south hemisphere that isquite compact to that seen in the north. It clearly shows
asymmetry in the north and south structure and clearlyoﬀ-set from the galactic center.
A giant gamma-ray structure was discovered by Fermi
in the energy range of 1-10 GeV (Dobler et al., 2010).
They reported that a feature emerges from the galactic
center and extends 50◦ north and south from the planeof the Milky Way, forming a dumbbell shape. This suggests an explosion from the galactic center. There areseveral reports that the X-ray feature is connected with

13

Fig. 9. SSC color map of the entire sky where red, green and blue correspond to the energy ranges of 0.6-1 keV, 1-2 keV and 2-4 keV.

the Fermi bubble. Kataoka, et al. (2013) performed observation using Suzaku. Since it has a relatively small
FOV, of 22 square, they observed several points aroundthe edge of the Fermi bubble, north and south. They
detected an intensity or temperature jump at the edge
location they selected. The spectrum reported is in thermal origin, rather than non-thermal origin.
Comparing the Fermi bubble map with that of theSSC, we see that there is not so clear correlation between
them. In the north hemisphere, the Fermi bubble corresponds to the relatively weak emission in X-ray. The
X-ray feature seems to surround the Fermi bubble, suggesting the formation of the shock front. In the south
hemisphere, however, the Fermi bubble surely includes
the X-ray bright region.
We have accumulated the spectrum from the extended
region covering the Fermi bubble shown in ﬁgure 10. We
ﬁnd that it can not be non-thermal origin. Although
we have to modify the abundance for Ne and Mg, it is athin thermal origin. This result is consistent with that by
the Suzaku observation (Kataoka et al., 2013), although
Suzaku observed a very small region.
Then, we divide the X-ray structure into several segments; inside the Fermi bubble, outside the Fermi bubble, north, south according by latitude and longitude.
We ﬁnd that the plasma temperature is almost constant,
0.31 keV whole through the region. The absorption feature, NH , is a reliable distance indicator in X-ray. However, we can not determine NH due to the fact thatour data are insensitive below 0.6 keV. According to the
ROSAT data, there is an intensity variation at low energy due to the local hot bubble or the solar wind charge

Fig. 10. SSC spectrum around the Fermi bubble region

exchange or whatever. It will be quite diﬃcult to estimate the value of NH even if we are sensitive to low
energy. We need higher energy resolution below 0.6 keV.
Figure 11 shows a simple sphere projection using
JUDO2 (http://www.darts.isas.jaxa.jp/astro/judo2/).
The ﬁgure center is selected by the center of the X-ray
structure so that almost all the extended emission canbe easily seen here. It is about 20◦ away from thegalactic center and not on the galactic plane. There aremany galactic sources along the galactic plane where
they contaminate the extended emission. Apart fromthem, we see that the soft X-ray emission is well withina circle with a radius of 50◦ . The north polar spur is14using good energy resolution will be able to determine
the value of NH to decide whether the X-ray structureis around the galactic center or is very close to us. We
expect that the near future observation with eROSITA(http://www.mpe.mpg.de/eROSITA) will solve thisproblem.

The SSC operation is supported by many people in
the MAXI team as well as students who work outside
the team at present. We greatly appreciate all the eﬀorts
they have done so far.

Fig. 11. Diﬀuse structure around the Fermi bubble region

relatively bright in the north-east edge. Other thanthat, it reminds us a center-ﬁlled SNR with a large
scale. Pink color spots in ﬁgure 11 indicate the FOVof the Suzaku observations performed so far, most ofthem are for the speciﬁc point sources. Very few are
for the study of the extended structure (Kataoka et al.,2013). They can cover much less than 1% of the X-ray
large structure. We have measured the average emission
integral (EI) of the large structure along the line ofsight. We ﬁnd EI 1 ∼ 4 × 10−2 cm−6 pc with which
we can estimate the total mass (Mt ) and the thermalenergy (Eth ) as below, where D is the distance to thesource and kT is the temperature.
Mt∼

6 × 104 (

D 2.5
) M
1 kpc

Eth

∼

4 × 1052

kT
D 2.5
(
) erg
0.3 keV 1 kpc

The X-ray structure showing circular shape is oﬀ-set
from the galactic center while the Fermi bubble is symmetry on the galactic center. In the northern hemisphere, the X-ray bright region surrounds the Fermi bubble while in the southern hemisphere, the X-ray brightregion is surrounded by the Fermi bubble. This suggeststhat they are accidentally seen in the same line of sight.
If the X-ray structure originates from the galactic center,
Mt is an order of 107 M and Eth is an order of 1055 erg.
This requires us a very strong AGN activity in ourgalaxy. If the source is well within one kpc, these two values become plausible, supplied by hypernova or a seriesof supernova. The precise distance determination is the
key to solve this problem. The observation below 0.6 keV

References
Amati, L., et al., 2008, MNRAS, 391, 577
Cash, W., et al. 1980 ApJ., 238, L71
Dobler, G., 2010, ApJ., 717, 825
Freyberg, M. J., et al., 1999, HXRA conf. 278
Imatani, R., et al. 2016 PASJ., 68, S19
Kataoka, J., et al. 2013 ApJ., 779, 57
Kimura, M., et al. 2013 PASJ., 65-1, 14
Marshall, F. E. et al. 2011 ApJ., 727, 132
Matsuoka, M., et al., 2009 PASJ., 61, 999
Mihara, T., et al., 2011 PASJ., 63S, 623
Miyata, E., et al., 1994 PASJ., 46L, 101
Morii, M., et al. 2013 ApJ., 779, 118
Romano, P., et al. 2006 A&A, 456, 917
Tomida, H., et al. 2011 PASJ., 63, 397
Tomida, H., et al. 2016 PASJ., 68, S32
Tsunemi, H., et al. 2010 PASJ., 62, 1371
Voges, W., et al. 1999 A& A, 349, 389
Yonetoku, E., et al., 2010 PASJ, 62, 1495

15

Discovery of 17 X-ray Transients with MAXI/GSC and Their Nature
Hitoshi Negoro, and MAXI team
Department of Physics, Nihon University
1-8 Kanda-Surugadai, Chiyoda-ku, Tokyo, Japan
E-mail(HN): negoro.hitoshi@nihon-u.ac.jp
Abstract
MAXI newly discovered 15 uncatalogued, and two unidentiﬁed X-ray transients for seven years. From
the MAXI and followup observations, especially thanks to rapid Swift/XRT followup observations, thosetransients were found to be six black hole candidates, six neutron stars including one pulsar, and one
white dwarf. The possible nature of the remaining four sources, especially MAXI J1957+032 and MAXI
J1501−026 recently discovered, are discussed. Characteristics and statistics of these X-ray transients arereviewed focusing on the discovery rate of the black hole candidates and the spatial distribution of thesetransients in our Galaxy. The peak ﬂuxes of the MAXI transients except for two black hole candidates,
MAXI J1659−152 and MAXI J1910−057, are relatively low and about 100 mCrab or less. This implies
that MAXI has observed distant areas in our Galaxy as was expected before the launch. On the other
hand, a high discovery rate of transients at high galactic latitudes suggests the presence of a large numberof low mass X-ray binaries wandering in the halo.
Key words: X-rays: transients, X-ray novae — binaries: black holes, neutron stars, white dwarfs

1.

Introduction

One of the characteristic features of X-ray binaries istransient behavior. They show various time variations,e.g., X-ray bursts, outbursts, and state transitions. The
number of the observed outbursts is a fundamental key totest stellar and binary evolution theories. In our Galaxy,for instance, 0.1–1 billion stellar black holes are expected
to be present theoretically, but only about 60 black hole
X-ray binaries (here after BHCs)1 have been discovered(e.g., Corral-Santana et al. 2016). And, more than 90%
of BHCs are transients, showing outbursts typically lastmore than 100 days (McClintock & Remillard 2006).
Thus, searching transient BHCs and neutron star binaries is important not only to study physical propertiesof those objects, but also to conﬁrm the evolution theories. MAXI, Monitor of All-sky X-ray Image (Matsuoka
et al. 2009), with two very wide ﬁelds of view (∼ 150◦ ×3◦each) is a mission focusing on the discovery of such transients rather than short-term rapid variability such as
X-ray bursts and gamma-ray bursts.
Of course, MAXI also detected a number of bursts*1 These black holes include dynamically mass-determined black
holes and black hole candidates showing very similar observational properties. Here, those black hole binaries are simply
denoted as black hole candidates (BHCs) because no direct evidence has not been obtained, yet, from both types of objects.

and ﬂares from various objects as was originally expected(Negoro et al. 2016a, hereafter NKS). Here, I summarize
the nature of 17 new X-ray transients (mostly novae)
discovered with GSC (Mihara et al. 2011; Sugizaki et al.
2011) for 7 years, and discuss the MAXI capability to
discover such transients and the spatial distribution of
the MAXI transients in our Galaxy.
2. 17 X-ray Transients Discovered with MAXI/GSC
MAXI/GSC solely or independently discovered 17 X-ray
transients in terms of the MAXI/GSC Nova-Alert System (NKS). The locations of these transients on galactic
coordinates are shown in ﬁgure 1. Observational properties of the 15 X-ray novae discovered in the ﬁrst 68
months are described in NKS. After that, MAXI newly
discovered a short X-ray transient MAXI J1957+032 on2015 May 11 (Negoro et al. 2015a), and a short and soft
X-ray transient MAXI J1501−026 on 2015 August 26(Nakahira et al. 2015).
2.1. Deﬁnition of MAXI transient sources
First of all, I decribe the deﬁnition of ”new MAXI transients” used here. The MAXI objects should be uncatalogued objects. Not only gamma-ray bursts (GRBs)
but also quiescent or faint known objects are not included even if an outburst from the source is observed
for the ﬁrst time. For instance, MAXI J0911−655/Swift16

Fig. 1. 17 X-ray transients MAXI discovered. Six black hole candidates shown in bold italic (pink), six neutron star binaries in italic (green),
one white dwarf binary in italic (red), and four unknowns. The background image was obtained by ﬁve years GSC observations and three
years SSC observations. Therefore, the new sources are not always visible in this map.

J0911.9−6452, of which an outburst was recognized forthe ﬁrst time in 2016 February, is a member of the globular cluster NGC 2808 (Serino et al. 2016). The source
has been identiﬁed as Chandra source No. 15 (Homan etal. 2016, also see Servillat et al. 2008). Thus, the source
was a known catalogued source, and not included in the
MAXI sources, here.
The identiﬁcation by other observatories is, therefore, very important. Of the 17 new transients, MAXI
J1619−383 (Serino et al. 2014) and MAXI J1501−026
have not been identiﬁed by other observatories, yet. Sun
constraints did now allow other observatories to observe
MAXI J1619−383 in outburst (NKS). In that case, radio observations are likely a powerful tool to localize theposition and to know the nature of the source from, forinstance, the radio and X-ray luminosity diagram (Corbel et al. 2013; Corbel 2017). From the X-ray spectrum,MAXI J1619−383 is likely to be a low-mass X-ray binary
hosting a neutron star (NKS). For MAXI J1501−026, see§2.3.
On the other hand, the following transients are excluded in the list even if Swift/XRT successfully caught
the (uncatalogued) sources: short hard transients, whichmight be associated with GRBs, and short (soft) transients detected only at a single (or a few) scan transit(s), having not enough statistics or information to
restrict the nature of the source. A former example
is MAXI J0636+146. MAXI J0636+146 was detectedsigniﬁcantly at, at least, two scan transits (Negoro etal. 2016b). Swift/XRT successfully observed the source,
but the nature of the source is still unclear (Kennea etal. 2016b; Kennea 2017). While, the latter soft transients includes MAXI Unidentiﬁed Short Soft Transients(MUSSUT) discussed by Mihara (2017).
In this sense, the deﬁnition is somewhat ambiguous.
Some may be excluded or included in future.
2.2. Category of the new transients
Currently, the 17 new transients are categorized into six
BHCs, six neutron star binaries, one white dwarf binary, and four unknowns (ﬁgure 1). MAXI discovered
the most BHCs, six, of 14 BHCs found after the MAXI
launch, though none of their mass is not dynamically determined. The observed properties of the six BHCs are
summarized in Negoro & MAXI team (2014) and NKS.
Six sources were found to be neutron star binaries from
regular pulses or X-ray bursts (NKS). MAXI J1932+091
with a probable Be companion (Itoh et al. 2014) is acandidate of a super-giant fast X-ray transient (NKS;
Sakamaki & Negoro 2017), and spectral properties of
MAXI J1619−383 are consistent with that of neutronstar LMXBs (NKS). MAXI J0158+744 is a quite rare,
rapidly fading super-soft-source on the outskirts of the
Small Magellanic Cloud (SMC) (Morii et al. 2013).
2.3. Nature of MAXI J1957+032 and MAXI J1501−026
What is the nature of the new transients of MAXIJ1957+032 and MAXI J1501−026? MAXI J1957+032
was discovered with MAXI/GSC on 2015 May 11 (Negoro et al. 2015a). INTEGRAL also independently discovered the source (Cherepashchuk et al. 2015). Thus,
it is known as IGR J19566+0326. As shown in ﬁgure
2, MAXI J1957+032 exhibited short X-ray outbursts

17

0.2

0.15

3. Discovery Rate of Black Hole Candidates
15/05/11

16/01/07

16/09/29(#7504)(#8529)(#9565)

15/10/06

Before the MAXI launch, in 2008, I optimistically estimated the discovery rate of BHCs with MAXI/GSC to
be 7.5–15 BHCs/yr, which was 5–10 times higher than
before (1.5 yr−1 ), taking into account the extension ofobservable regions, the star distribution in our Galaxy,and frequent scanning of the galactic center region (Negoro 2009).

Photons cmí2 sí1(#8143)

0.1

0.05

0

5.72×104

3.1. Expected and observed rates

5.74×104
MJD

5.76×104

Fig. 2. MAXI/GSC 2-20 keV light curve of MAXI J1957+032. The
numbers below the dates (YY/MM/DD) are reference numbers of
The Astronomer’s Telegram. 6 hr-bin data are used, and rebinned
up to 5 days as the detection signiﬁcance should be more than 3sigma level.

four times (2nd: Sugimoto et al. 2015; Kennea et al.
2015, 3rd: Tanaka et al. 2015, and 4th: Negoro et al.
2016c; Kennea et al. 2016a).
X-ray ﬂuxes of MAXI J1957+032 rapidly decreased
to below the detection limit of MAXI/GSC typically inless than one day. As a result, the nature of the source
had been long in mystery. Extensive Swift/XRT followup observations for the 4th outburst (Kennea et al.
2016a), however, revealed a clear correlation between theﬂux and the power-law photon index, strongly suggestingthat the source is a neutron star LMXB (Kennea 2017).
On the other hand, MAXI J0501−026 is a quite mysterious object. The source was detected at 440 mCrab
only at a single scan transit on 2015 August 26, and its
spectrum was extremely soft, e.g., kTBB = 0.53 ± 0.07
keV for a blackbody ﬁt (Nakahira et al. 2015). These
properties are reminiscent of the SSS source MAXI
J0158+744. However, no X-ray (Swift/XRT; J. Kennea
private comm.) and optical (KWFC; Morokuma et al.
2015) counterpart was found.
If the object really emitted a blackbody spectrum,from the normalization of the spectral ﬁt, we can estimate the source distance to be ∼ 1 kpc for a neutron starwith a radius of 11 km, or ∼ 640 kpc for a white dwarf
with a radius of 7,000 km. Thus, the nature of the sourceis quite puzzling. I exclude this source in the following
discussion though it is a quite high galactic source.
In summary, the nature of the 17 new X-ray transientsare very likely six BHCs, nine neutron star binaries, one
white dwarf binary, and one (quite) unknown.

The expected rate is unexpectedly higher than the current rate of 2.0 yr−1 not only by MAXI but also by Swiftand INTEGRAL (ﬁgure 3). One reason of the discrepancy is that unfortunately no bright (> 1 Crab) novae
have appeared for more than ten years, which decreases
the rate by 1–2. Furthermore, no observatories havefound even faint (< 0.1 Crab) BHCs since the discovery of IGR J17454-2919 at the end of September, 2014
(Chenevez et al. 2014). This also reduces the rate by∼ 1. Another reason is that the past discovery rate,
which was estimated from 29 novae2 from 1987 to 2007,
was already increased by, for instance, deep observationsof the galactic plane with RXTE/PCA (not ASM), INTEGRAL, and Swift.
Furthermore, more importantly, new outbursts from asimilar number of known BHCs to new ones were also
detected (see Table 2 in NKS for MAXI detections until
2015 March). For instance, XTE J1856+053 (Suzuki etal. 2015), XTE J1908+094 (Krimm et al. 2013; Negoro
et al. 2013), GS 1354−645 (Miller et al. 2015) and GS
2023+338/V404 Cygni (Barthelmy, et al. 2015; Negoro
et al. 2015b; also see Rodriguez et al. 2015) showed renewed activity for the ﬁrst time in 8, 10, 18, and 26 years,
respectively. This potentially reduces the chance of discovery for later missions. A recurrence rate of outbursts
reduced from such detections is, however, very important to estimate the numbers of BH binaries and BHs in
our Galaxy.
Certainly, it is diﬃcult for MAXI to discover outbursts close to bright sources due to moderate spatial resolution as an all-sky monitor. INTEGRAL and
Swift, however, detected such sources, for instance, Swift
J174510.8−262411 (Cummings et al. 2012) near GX 3+1
and IGR J17454−2919 (Chenevez et al. 2014) near thegalactic center. Thus, this is unlikely to help the large
diﬀerence so much.
From the above discussion, the observed rate may increase to ∼ 5 if the X-ray sky is just like before and if
we look for faint sources on longer time scales (see §5).
*2 Later, some were found to be BHCs, and the current numberof BHCs newly discovered in this period is 34, which yields therate of 1.7 yr−1

18Fig. 3. Historical numbers of black hole candidates discovered by various missions. The numbers shown here are slightly diﬀerent from thoseshown by Corral-Santana et al. (2016). This is due to the diﬀerence in the deﬁnition of BHCs.

3.2. Distant black hole candidates?
MAXI indeed detected faint X-ray novae as was
expected.
The peak ﬂuxes of MAXI J1543−564,
J1836−194, J1305−704, and J1828−249 were approximately 100 mCrab, which were much smaller than thoseof previous bright novae by two orders of magnitudes(NKS). The peak luminosities of BH novae are distributed around 10% of the Eddington luminosity, LEdd
(Chen et al. 1997; Dunn et al. 2010).
On the other hand, of the above four MAXI transients, MAXI J1543+564 (Kennea et al. 2011; Stileet al. 2012), J1305−704 (Morihana et al. 2013), andJ1828−249 (Tomsick & Corbel 2014) also exhibited thesoft-to-hard (or soft-to-intermediate) state transitions.
The soft-to-hard state transitions are known to occur atabout 1–4% LEdd for nine sources including three NS
binaries (Maccarone 2003) and in the range of 0.5–10%
LEdd with the peak at 2–3% LEdd for 25 BHCs (Dunnet al. 2010, also see Masumitsu & Negoro 2017).
From these facts, these sources are thought to be relatively far from the solar system, ≥ 12 kpc, e.g., for MAXI
J1543+564 (Stile et al. 2012).
Thus, it can be said that MAXI is actually observing
distant parts of our Galaxy, and detecting distant galactic transients, though half GSCs have not been used.
4. Distribution of the MAXI Transients and Its Suggestion
4.1. High galactic latitudes sources
Let me go back to all the X-ray transients, but now excluding the quite unknown source MAXI 1501−026 andthe SMC source MAXI J0158−744. One of the notice-able features of the MAXI transients is that some of them
are located at relatively high galactic latitudes as shownin ﬁgure 1. One out of six BHCs, ﬁve or six of nine
NS binaries are at galactic latitudes of more than 8 deg.
This suggests the presence of a number of LMXBs in thegalactic halo. Note that MAXI J1932+091 is the only
high mass X-ray binary (candidate) in the MAXI transients. It is known that bright LMXBs are often in globular clusters (e.g., Predehl et al. 1991). But, we could
not ﬁnd any globular clusters in these source directions.
Figure 4 shows cumulative numbers of LMXBs
counted from the high galactic latitude. These numbers
are obtained from the 4th LMXB catalogue by Liu et al.
(2007), which contains 187 LMXBs in our Galaxy, theLMC (Large Magellanic Cloud), and the SMC. In thehistogram, globular cluster members, the nearby (∼ 200pc) LMXB 4U 1700+24, LMC and SMC members are
all excluded for this speciﬁc purpose3 .
Only 6% (10 of 175) known LMXBs are located at
galactic latitudes |b| > 15.8 deg, but 20% (3 of 15) MAXItransients were found at the high latitudes. The signiﬁcance of the large discrepancy between these ratios is low.
But the signiﬁcance becomes more signiﬁcant if lower latitude sources are included; 9% (15/175) vs. 33% (5/15)
at |b| > 12.4 deg, and 18% (31/175) vs. 40% (6/15) at|b| > 7.6 deg. Thus, MAXI detected high galactic latitude sources 2–3 times larger than expected from the
distribution of known LMXBs, though the signiﬁcanceis not enough, yet.
*3 For 4U 1543−47, the galactic latitude of +5.4 deg is used, not+54. deg in the catalogue.

19Fig. 4. Logarithmically binned cumulative numbers of galactic LMXBs counted from the high galactic latitude from the LMXB catalogue by
Liu et al. (2007) and MAXI transients. The left side bin shows the number of |b| < 0.5 deg, exceptionally.

This suggests the presence of a large number of unknown/sleeping LMXBs at high galactic latitudes unless
MAXI happened to ﬁnd many transients at high latitudes. In that case, this also implies that MAXI failed todetect a much larger number of transient LMXBs (about
30 LMXBs) on or near the galactic plain because of thepoor spatial resolution of the MAXI cameras. If so, INTEGRAL or Swift (and RXTE) should detect a large
fraction of such sources during their galactic scan survey.
But, the discovery rate of such sources is not so high.
This might indicate that the transient LMXBs have adiﬀerent galactic distribution from the known LMXBs.
4.2. LMXBs wandering in the halo?
Figure 5 shows an edge-on view of our Galaxy and locations of the MAXI transient sources projected onto an
x-z plane, where the x–axis is taken to be alone the direction from the Solar system to the galactic center. Somedistances to the sources are taken from literatures. Someare estimated from luminosities at the state transitions
for BHCs as described in §3.2, and for MAXI J0556−332(Sugizaki et al. 2013). Some NS binaries can be also estimated from the peak ﬂuxes of X-ray bursts (NKS) as
reference to Kuulkers et al. (2003). The distances to
MAXI J1619−383, J1932+091, and J1957+032 are unknown, and assumed to be 4 kpc without any strong
reasons except for J1957+032 (see Kennea 2017).
It can be found from the ﬁgure that more than half
MAXI transients are not close to the solar system, but

far and above the galactic disk. Thus, it is possible thatMAXI observed runaway LMXBs on the analogy of runaway radio pulsars. To estimate the number of such
LMXBs, we need more samples and more quantitative
discussion for the detection or non detection for galactic plane transient sources. Measurement of the propermotion of the transients is, of course, very important as
was the case of XTE J1118+480 (Mirabel et al. 2001).
5. Discussion and Future Works
Finally, a strategy to ﬁnd fainter X-ray novae is described. Currently, the longest time bin to ﬁnd transient
events in the nova-alert system is 4 d bin, and the detection limit with the bin is about 8 mCrab (NKS). But,
if we use 16 or 32 d bin, then the detection limit canbe decreased to 5 mCrab or less. I have already conﬁrmed using archival data that previously undetectable
faint X-ray novae, e.g., IGR J17177−3656 and Swift
J1357.2−0933, and some tidal disruption events triggered the nova-alert system!
Furthermore, reanalyzing all the data for more thanseven years is also in preparation by K. Tanaka at Nihon university to obtain the complete sample of X-ray
transients.
The detection of the high galactic latitude sources
are entirely due to high sensitive, and long-term all-skysurvey with MAXI/GSC. Such detection may reveal thenature of a possible large number of previously unknown

20

J1659-152
Sun

J1647-227

J1619-383

J1910-057
J1543-584
J1305-704 J1828-249
121225A

J1932+091
J1957+032

J0556-332
Fig. 5. An edge-on view of our Galaxy, and projected distances of the MAXI transients on the plane. The sources in bold black font andgray font are BHCs and NS binaries, respectively. The sources to which the distances have been estimated are shown in italic font. The
distances to the other sources are assumed to be 4 kpc.

LMXBs wandering in the halo. More quantitative estimation and discussion will be reported elsewhere. In
this sense, too, further continuous observations by MAXIin future and the detection of transient events must bevery important not only for the stellar and binary evolution/formation theories, but also for previously unknown
LMXBs wandering in the halo.
This work was supported by JSPS KAKENHI Grant
Numbers 24540239 and 16K05301.
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Liu, Q. Z., van Paradijs, J., van den Heuvel, E. P. J.
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Masumitsu, T., Negoro, H. 2017, in this volume.
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5886

21

MAXI Data Archive Plan
K. Ebisawa1 , S. Nahahira1 , Y. Sugawara1 , S. Ueno1 , H. Tomida1 , T. Furusho1
T. Mihara2 , M. Sugizaki2 , M. Shidatsu2 , M. Serino2
H. Negoro3 , Y. Nakagawa4 , L. Angelini51Japan Aerospace Exploration Agency (JAXA), 2 RIKEN,
3
Nihon University, 4 JAMSTEC, 5 NASA/GSFC
E-mail(KE): ebisawa@isas.jaxa.jp
Abstract

We describe the MAXI data archive plan, including the long-term archive after the MAXI project isﬁnished. One one-hand, while MAXI is in operation, daily light-curves and other data products are createdat and released from RIKEN. On the other hand, DARTS, which is JAXA’s multidisciplinary space-sciencearchive, is going to maintain the high-level scientiﬁc products from MAXI, even after the MAXI project isﬁnished. We are developing the MAXI archive to be released from DARTS, with which users can timely
access the full MAXI datasets, perform the MAXI data analysis easily, just like standard X-ray pointing
satellites. The MAXI archive will be also released from HEASARC at NASA, such that the analysis tools
and calibration ﬁles are included in the standard HEASOFT (ftools) and CALDB packages, respectively.
In addition, DARTS will provide convenient web-application tools with which users can display MAXI
images for various time-periods, together with other astronomical images/catalogs, and quickly produce
spectra, light-curves and responses via standard web-browsers.
Key words: archives : DARTS — HEASARC

1.

JAXA’s Science Archive: DARTSDARTS1 (Data ARchive and Transmission System) is
JAXA’s space science data archive, which primarilyarchives high-level data products obtained by JAXA’s
space science missions for not only astrophysics (X-rays,
radio, infrared), but also solar physics, solar-terrestrial
physics, lunar and planetary sciences. Microgravity science data and other experimental/observation data obtained by International Space Station are also archived
at DARTS.
DARTS is operated in ISAS (Institute of Space andAstronautical Science), located in Sagamihara, Kanagawa. While MAXI is operated at TKSC (TsukubaSpace Center) and major data processing is performed
at RIKEN, we have started MAXI data archive development at DARTS in 2013, envisioning the long-termarchive even after the MAXI project is ﬁnished. By
having MAXI archive at DARTS, long-term preservation
and accessibility of the data are secured, and interoperability with other astronomical database is enhanced.

*1 http://darts.isas.jaxa.jp

2. MAXI Archive at RIKEN
2.1. Science Products
While MAXI is in operation, data are regularly processed
at RIKEN, and daily all-sky images, light-curves of major sources and other scientiﬁc products are created andreleased from the RIKEN MAXI site2 . Users are free to
use these products for scientiﬁc research and publication,
but if they are interested in other targets not in the listor particular time-periods not adopted in the standard
processing, they need to use the “on-demand” analysis
system (below).
2.2. On-demand Analysis System
The MAXI on-demand analysis system3 (Nakahira et al.
2013 ) is developed so that users can carry out their own
customized data analysis using only web-browsers.
Users specify target position, time-period, energy
band and bin-width of the light-curve on web-browser,
then “submit” the job and wait to receive image, lightcurve, spectrum and response in the standard FITS format.
Users can carry out customized data analysis using
the on-demand system, but there are several limitations.
*2 http://maxi.riken.jp*3 http://maxi.riken.jp/mxondem

22Fig. 1. Top: An example of JUDO2 display, where MAXI GSC and Swift-BAT images are superposed (Galactic Center region). BAT catalog(green cross) and MAXI catalog (blue triangle) sources are also superposed. Bottom: Galactic plane around Vela region by SSC (left),
GSC (right), and SSC and GSC superposed(center). Some hard X-ray sources (such as Cen A, located around north-east corner) clearlyseen with GSC are hardly seen with SSC, and vice versa (such as Vela supernova remnant, on the plane toward west).

First, the system is running on a single machine, so theload is heavy when there are several users. Also, it is
not practical to carry out analysis for long periods oftime (it takes ∼400 minutes to analyze a single source
for the 7 year period) or many targets simultaneously. In
addition, the system uses proprietary data and software
maintained practically by a single developer, so the longterm maintenance is a challenge.
Consequently, the on-demand system will be ported to
UDON2 at DARTS (section 3.2.), envisioning long-term
operation.
3. Web-application Tools at DARTS
In principle, DARTS provides high-level scientiﬁc data
products to download, so that users can carry out
data analysis in their own sites. In addition, DARTS
has various web-application services to facilitate quickbrowse/analysis of the data archived at DARTS.
3.1. Interactive image browsing system : JUDO2
JUDO24 is second generation of JAXA Universe Data
Oriented. JUDO2 adopts Aladin Lite (Boch and Fernique 2014), with which various all-sky data archived
at DARTS and other sites can be displayed on a webbrowser. Users can easily manipulate images (move,*4 http://darts.isas.jaxa.jp/astro/judo2

zoom-in/out), and superpose two images by adjustingtransparency. It is also possible to superpose various
marks (catalog sources etc.) and drawings (instrumentﬁeld-of-views, constellation boundaries, etc.) on the displayed image.
In Figure 1 (top), we show an example of displaying
MAXI GSC image and Swift BAT image superposed oneach other. In addition, MAXI catalog and Swift-BAT
catalog sources are indicated. It is particularly useful tocompare two instruments having diﬀerent spatial resolutions, since often MAXI sources are confused, which areresolved by SWIFT-BAT. Also, it is useful to display two
images simultaneously taken in diﬀerent energy bands,
such as MAXI GSC and SSC images (Figure 1, bottom).
We also provide MAXI all sky images integrated for,not only the entire mission life, but also each year,
month, week and day, so that evolution of transientsand variable sources are easily seen. The all sky image
data used in JUDO2/Aladin Lite are in the HiPS format(Fernique et al. 2015 ) and public5 , so that any Aladin
or Aladin Lite users can display MAXI images on theirlocal machines.

*5 http://darts.isas.jsxa.jp/pub/juod2/HiPS

23

Fig. 2. An example of link from JUDO2 to UDON2, where MAXI GSC and Suzaku images of 4U1957+11 are displayed on JUDO2 (left).
While the source is in the ﬁeld of view of JUDO2, if the UDON2 button (upper-right) is clicked, a new windows is open (center) fromwhich users can launch Suzaku UDON2 (right top) or MAXI on-demand (right bottom) for further quick-analysis of the data. Until thetime when MAXI on-demand is implemented in UDON2, the MAXI link in JUDO2 points RIKEN on-demand instead of UDON2.

3.2. Quick Analysis Tool : UDON2
UDON26 is the second version of Universe via DARTS
ON-line, with which users can quickly analyze astronomical data archived at DARTS using web-browser. As of
February 2017, ASCA and Suzaku data are available on
UDON2, and Hitomi data is being implemented (public
release in September 2017). In UDON2, users choose anobservation to display the image, select a region on theimage to extract spectrum and light curve. Once spectrum and light curve are displayed, users can change
scale and zoom-in to scrutinize minute features. Since
UDON2 does not depend on proprietary data and software (besides minor web-interface/display parts), longterm maintenance is almost guaranteed. Once MAXI
data and analysis software are made public, the ondemand system in RIKEN will be ported to UDON2 toensure long-term operation.
Also, it is easy to jump from JUDO2 to UDON2 (Figure 2). In this manner, users can choose a favorite
source/region on the sky, compare images taken by different instruments (using JUDO2), and closely inspect
the spectrum and light curve taken by each instrument(using UDON2). This can be done only using a standard
web-browser.
4. Long-term MAXI Archive at DARTS and HEASARC
4.1. Archive Data Structure
We are developing the full MAXI data archive, where
all the MAXI events are stored in the standard event
FITS format, commonly used in high energy astrophysics*6 http://darts.isas.jaxa.jp/astro/udon2/udon2.html(Nakagawa et al. 2016). The sky is divided into 768
HEALPix regions, and daily events from each region is
kept in a single event FITS ﬁle. In this manner, users cananalyze MAXI data just like standard pointing satellite
data. Detail of the archive directory structure is foundin Furusho et al. (2017).
4.2. Archive Data Flow
Pipe-line processing of the MAXI data is carried out at
TKSC (Figure 3). There are several routes for MAXI
telemetry data to arrive TKSC from ISS, so that thepacket loss and/or deﬁciency are not uncommon. Thus,
pipe-line processing is continuously running at TKSC,
to process not only the latest observation sequences,
but also previous sequences that have been already processed. All the processed data are transferred to DARTS,
and the old data previously processed are overwritten by
the newly processed data corresponding to the same observation. In this manner, the data archive will get closer
to completion with time. The data are publicly released
from DARTS and permanently archived. The data aremirrored from DARTS to HEASARC, so that users candownload the identical datasets either from DARTS or
HEASARC.
4.3. Archive Data Download
MAXI archive data cover almost the entire sky for overseven years. A straightforward way to fully explore the
MAXI archive is to download (mirror) the entire archive
either from DARTS or HEASARC to users’ local sites(top horizontal line in Figure 4.). This is not infeasible,
with the current total archive size ∼1.4 TB. In fact, we24

Fig. 3. MAXI archive data ﬂow. Pipe-line processing is performed at TKSC, and the processed data are continuously sent to DARTS. These
processed data are mirrored to HEASARC. Archive users may download the data either from DARTS or HEASARC.

Fig. 4. MAXI archive data analysis ﬂow. Usually, archive users download a part of the archive data either from DARTS or HEASARC,
using “mxdownload” or speciﬁc download tools. Alternatively, users may mirror the entire MAXI archive on their own sites. Users run“mxproduct” script either on the download data or the local archive to produce science products.

will provide wget scripts to download the entire MAXI
archive from DARTS.
However, many users are interested in only particulartargets or regions on the sky for particular observation
periods. So, we will provide mxdownload script, withwhich users can download a part of the MAXI archivefrom DARTS for a particular sky region and observation
period. The downloaded dataset has the same directory
structure as the original MAXI archive, and considered
to be a subset of the archive.
4.4. Archive Data Analysis
MAXI data analysis can be carried out by combining
MAXI speciﬁc ftools and the standard ftools includedin the HEASOFT package. The number of MAXI speciﬁc ftools is only nine, one of which is a perl script,
mxproduct, to run the other eight ftools written in C.
mxproduct runs both on the entire archive or a partof the archive downloaded by mxdownload. In practice,users need to run only mxproduct to carry out analysis, where many options are available for data selection.
Explanation of these ftools are found in Furusho et al.
(2017). Outputs of mxproduct are image, light-curve,
spectrum and response in the standard FITS format.
Users can analyze these science products using popular
HEASOFT tools (spectral analysis with xspec, display

light-curve with lcurve etc.).
5. Schedule
As of February 2017, the MAXI archive system is being
tested actively by the MAXI team, and the eﬀort to implement the system to HEASARC is just started. We
hope to publicly release the archive data and software in
early 2017. After the release, the data are maintained
by DARTS (mirrored to HEASARC), and the softwares
and calibration ﬁles are maintained by HEASARC (incooperation with the MAXI team).
After the MAXI project is over in future, all thedata products created and released at RIKEN will beported to DARTS and permanently archived, together
with JAXA’s other science data archives.
In this maner, we expect the MAXI data archive keep
producing scientiﬁc outputs continuously for long future.
References
Boch, T. and Fernique, P. 2014, ASPC, 485, 277
Fernique, P. et al. 2015, A. and Ap. 578, 114
Furusho, T. et al. 2017, in this volume
Nakagawa, Y. et al. 2016, Journal of Space Science
Informatics Japan: Vol.5, p.71
Nakahira, S. et al. 2013, Journal of Space Science
Informatics Japan: Vol.2, p.29

25

26

27

The MAXI/GSC catalog in the low Galactic-latitude sky
Takafumi Hori,1 Yoshihiro Ueda,1 Megumi Shidatsu,2Taiki Kawamuro,1 Atsushi Tanimoto,1 Saeko Oda,1 and the MAXI team1

Department of Astronomy, Kyoto University, Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan
2
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(TH): hori@kusastro.kyoto-u.ac.jp
Abstract

Complete X-ray source catalogs in our Galaxy are the basis to study the demography of Galactic X-ray
populations. Since MAXI/Gas Slit Camera (GSC) covers the 2–10 keV band with the highest sensitivity
as an all-sky mission, its Galactic source catalog will be complementary to those of hard X-ray (>10 keV)
missions. Here we present the latest X-ray source catalog of MAXI/GSC in the low Galactic-latitude sky(|b| < 10◦ ). We utilize the same image ﬁtting method as that in Hiroi et al. (2013), who constructed
the 37-month MAXI/GSC catalog at high Galactic latitudes (|b| > 10◦ ). To overcome source confusion incrowded regions, we have accurately calibrated the position-dependent shape of the point spread functionof MAXI/ GSC by utilizing the onboard data. We also take into account the Galactic ridge X-ray emission
in the background, utilizing the model by Revnivtsev et al. (2006). With a preliminary analysis, we detect
149 sources with signiﬁcance above 8σ, among which about 71% have counterparts in the Swift/BAT
70-month catalog.
Key words: workshop: surveys — X-rays: catalog

1. Introduction
MAXI/Gas Slit Camera (GSC) has achieved the best
sensitivity in the energy band of 2 − 10 keV among all
past or on-going all-sky X-ray missions. All-sky X-ray
source catalogs of MAXI/GSC are complementary toother ones covering diﬀerent energy bands, such as thoseof ROSAT (< 2 keV), Swift (> 14 keV), and INTEGRAL (> 15 keV). The 37 month MAXI/GSC catalog
in the high Galactic-latitude sky (|b| > 10◦ ) is publishedby Hiroi et al. (2013). At low Galactic latitudes, source
confusion becomes more signiﬁcant because the Galactic plane is crowded by bright X-ray sources. Also, weneed to take into account the the Galactic ridge X-ray
emissions (GRXEs) in modeling the background. In thispaper, we report the current results of the MAXI/GSC
catalog in the low Galactic-latitude sky.
2. Data Reduction and Analysis
We utilize the ﬁrst 37-month MAXI/GSC data observed
from 2009 September 23 to 2012 October 15. The Galactic center region (l < 30◦ , 330◦ < l, and |b| < 5◦ ) is
excluded, where source confusion is too serious. Data
reduction is performed with the same procedures described in Hiroi et al. (2013). We use our latest model ofthe non X-ray background (NXB) plus the cosmic X-ray

background (CXB), which is detailed in Shidatsu et al.
(2017, this conference).
In addition to the NXB and CXB, we consider the
GRXE, whose surface density rapidly increases toward
the Galactic plane (b = 0). We model the GRXE proﬁle with the ellipsoid model by Revnivtsev et al. (2006),
which consists of the Galactic disk and bulge/bar components.
3.

Current Results

To determine the signiﬁcance and ﬂux of each X-ray
source, we perform Poisson maximum-likelihood ﬁt to
the observed image with a model consisting of source
point spread functions (PSFs) and the background(NXB+CXB+GRXE). We ﬁrst only consider the sourcesof the Swift/BAT 70-month catalog by ﬁxing their positions. Then, from the residual image, we search forsources detected only with MAXI/GSC.
As a result, we detect 149 sources with signiﬁcanceabove 8 σ at |b| < 10◦ , among which 107 (71%) are in the
Swift/BAT 70-month catalog. The number of sourcesin each category is summarized in Table 1. Figure 1
shows the locations of the detected sources in Galactic
coordinates, including the high Galactic-latitude sourcesin Hiroi et al. (2013). Figure 2 shows the log N -log

28

Number of Sources
1
11
19
6
11
6
45
(25)
(6)
4
4
42

S relation in the low Galactic-latitude (|b| < 10◦ ) sky
outside the Galactic center region, where the number
density N of all sources with 4–10 keV ﬂuxes above S isplotted. The slope is close to unity.
4.

10í3

Category
Galaxies
Galaxy clusters
Seyfert galaxies
Quasars
Cataclysmic variables
Stars
X-ray binaries(Neutron star)(Black Hole)
Pulsars
Supernova remnants
Unmatched

N (>S) [degí2]

Table 1. Categories of Cataloged Sources

0.01

0.1

Fig. 1. The locations of all MAXI sources in Galactic coordinates, including those of Hiroi et al. (2013) obtained from the 37 month data
at high Galactic latitudes. Each radius is proportional to the logarithm of the ﬂux. The Diﬀerent colors correspond to diﬀerent types;
galaxies (cyan), galaxy clusters (green), Seyfert galaxies (blue), quasars (purple), cataclysmic variables (purple), stars (orange), X-ray
binaries (red), pulsars (black), supernova remnants (pink) and X-ray sources without counterpart in Swift/BAT 70-month catalog (Gray).

10í12

10í11
10í10
10í9
10í8
S (4í10 keV) [ergs cmí2 sí1]

Fig. 2. The log N –log S relationof the MAXI/GSC sources in the
Galactic plane region.

PSFs as accurately as possible. We are re-calibrating
the PSF proﬁles by referring to the onboard data indetail. This new PSF database will be utilized for all
future works of catalog production (including that at
high Galactic-latitudes) and light curve extraction from
MAXI/GSC data.

Future WorkWe are working to produce a MAXI/GSC low Galacticlatitude catalog utilizing the 7-year data (Hori et al.
2017, in preparation), which will update the results presented here. To overcome the diﬃculties in detecting
faint objects around bright sources, we must model the

10í7

References
Revnivtsev et al. 2006, A&A, 452, 169
Hiroi et al. 2013, ApJS, 207, 36

29

MAXI/GSC 7-year all-sky map and studies of large-scale
diﬀuse structures
Megumi Shidatsu,1 Yoshihiro Ueda,2 Taiki Kawamuro,2Takafumi Hori,2 Satoshi Nakahira,3 and the MAXI team
1

MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
Department of Astronomy, Kyoto University, Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan
3
ISS Science Project Oﬃce, Institute of Space and Astronautical Science (ISAS), Japan Aerospace
Exploration Agency (JAXA), 2-1-1 Sengen, Tsukuba, Ibaraki 305-8505, Japan
E-mail(MS): megumi.shidatsu@riken.jp
2

Abstract
We present MAXI/GSC all-sky maps in the 3–4 keV, 4–10 keV, and 10–20 keV bands constructedfrom the ∼7-year data. To realize maximum sensitivity, we have developed a background model with highaccuracy, by using the rate of the coincident events between a signal anode and a veto anode, as a tracerof the intensity and energy distribution of the GSC background. The model has enabled us to obtain thehigh-quality all-sky maps and to analyze X-ray sources all over the sky. We report our studies of extended
structures, including the spatial distribution of the Galactic ridge X-ray background and constraints on
the strength of the Fermi bubbles in the X-ray band.
Key words: surveys — X-rays: diﬀuse background — Galaxy: structure

2. Data Reduction and Creation of All-sky Images
For data screening and the production of simulated background data, we basically followed the procedure described in Hiroi et al. (2013). The main diﬀerence between that work and ours is that we created the GSC
background model separately for the periods when theinternational space station (ISS) moves northward andsouthward. In Figure 1 we plot the correlations between
the background event rate and the “VC count rate”,
which is the rate of events coincidently detected by a signal anode and a vedo anode. As reported in Hiroi et al.

(2013), the background levels in terms of the VC countrate depend on whether or not the Russian spacecraft
Soyuz was attached to the ISS’s docking port close to
MAXI. The direction of the ISS motion also aﬀects the
GSC background: the background rate become higher
when the ISS travels from south to north than from northto south.
Acam 215

BG rate (cnts/sec)

1. Introduction
All-sky X-ray surveys with good sensitivity are important for studies of large-scale diﬀuse structures such as
the Galactic ridge emission and Fermi bubbles. The
MAXI (Matsuoka et al. 2009)/GSC (Mihara et al. 2011)
has achieved the highest sensitivity in the 4–10 keV band
than any other previous all-sky X-ray surveys (Hiroi etal. 2013) and thus it is suitable to study diﬀuse sources inthis energy range. In this article, we present MAXI/GSC
all-sky maps created by compiling the GSC data for almost the entire period of the MAXI’s operation, andshow some example of studies on diﬀuse X-ray sources
using the maps. These results will be reported in moredetail in a later paper (Shidatsu et al. in preparation).

10

Soyuz dock

Phase 3, north5

Soyuz undock
Phase 0, north

100

200

300
VC (cnts/s)

400

500

Fig. 1. Correlations between the VC count rates versus the 3–20 keV
background event rates in Acam 2. The black and red data (withhigher and lower background rates) show the correlations when
Soyuz is present and not present on the docking port close to
MAXI, respectively. The open circles and squares plot the datawhen the ISS travels northward and southward, respectively.

30

Fig. 2. Exposure-corrected, background subtracted MAXI/GSC all-sky images in 3–4 keV (red), 4–10 keV (green), and 10–20 keV (blue)
from 2009 August to 2015 October in Galactic coordinates. A color version of this ﬁgure is available online.

the 3–10 keV longitudinal proﬁle of the GRXE averaged
in |b| < 2◦ obtained from our all-sky map. We foundthat the proﬁle within |l| < 100◦ (where l is the Galactic
longitude) is consistent with that obtained from the 3–20
keV RXTE/PCA data (Revnivtsev et al. 2006). The intensity for |l| > 100◦ and correlations of the entire GRXE
proﬁle with the ridge emissions in other wavelengths areleft as future studies.
20|b| < 2 deg

surface brightness (1eí8 erg/cm2/s/str)

Considering these results, we constructed the GSC
background model separately for the periods when ISS
moved north and south and when Soyuz was docked
and undocked. We produced simulated background data
based on our model using the MAXI simulator (Eguchi etal. 2009). We found that our new background model improved the reproducibility of the GSC background proﬁleparticularly above ≈5 keV.
We created all-sky images of the real data and the simulated background data and the eﬀective exposure from2009 August to 2015 October. The exposure map was
produced in the same manner as Hiroi et al. (2013). Subtracting the background from the real image and dividing
the resultant image by the exposure map, we obtained
a background-subtracted, exposure-corrected all-sky Xray image. Figure 2 presents the all-sky X-ray maps in
3–4 keV, 4–10 keV, and 10–20 keV.

15

10

5

0

3. Studies of large-scale X-ray structures
Here we show examples of scientiﬁc results that canbe derived from our all-sky maps. First, we searched
the GSC all-sky maps for structures associated with the
Fermi bubbles. No signiﬁcant excess is detected, however, in and around the bubbles, after removing the data
for the Galactic plane (|b| < 12.5◦ , where b is the Galactic latitude) and point sources listed in the MAXI 7month catalog (Hiroi et al. 2011). We estimated theupper limit of the 3–10 keV intensity averaged for theentire region of each bubble to be ∼ 2 × 10−8 erg cm−2
s−1 str−1 .
Next, we investigated the spatial distribution of the
Galactic ridge X-ray emission (GRXE). Figure 3 presents

í100

0Galactic longitude (deg)100

Fig. 3. Longitudinal proﬁle of the GRXE obtained from the 3–10 keV
all-sky map.

References
Eguchi, S. et al. 2009, in Proceedings of the RIKEN
Symposium, Astrophysics with All-Sky X-Ray
Observations, ed. T. Mihara & M. Serino, 44
Hiroi, K. et al. 2011, PASJ, 63, 677
Hiroi, K. et al. 2013, ApJS, 207, 36
Matsuoka, M. et al. 2009, PASJ, 61, 999
Mihara, T., et al. 2011, PASJ, 63, 623
Revnivtsev et al. 2006, A&A, 452, 169

31

Operation and Calibration of MAXI GSC Xe Gas Counters on the ISS
for over 7 Years
Mutsumi Sugizaki1 on behalf of MAXI GSC Team1

RIKEN, MAXI Team, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(MS): sugizaki@riken.jp
Abstract

In-orbit operations of MAXI GSC Xe-gas counters on the on ISS for over 7 years, and their response
calibrations are presented. Passing through the operation for the 7 years since 2009, some gas counters
are found to expose performance degradation. Three out of the 12 gas counters are considered to havea broken anode, suﬀering repeated particle irradiation and discharge. Another gas counter continuously
exhibit a gain increase since 2013 June 10, suggesting a gas leak with an exponential decay time of ∼ 1.5
year. We operate these counters carefully using their alive components.
Key words: instrumentation: detectors

1. Overview of MAXI GSC
GSC (Gas Slit Camera; Mihara et al. 2011) is the main
onboard instrument that consists of 12 Xe-gas proportional counters optimized for the 2–30 keV energy bandand one-dimensional slat collimators. With the onedimensional wide ﬁelds of views (160◦ × 3◦ ), GSC scans
the whole sky every ISS orbital cycle (∼ 92 minutes).
The Xe-gas counters employ resistive carbon wires fortheir photon-detector anodes to achieve the position sensitivity. Each detector unit consists of two counterswhich work redundantly.
2. Status of in-orbit operation and performance
2.1. Daily sky coverage
Figure 1 shows an all-sky X-ray image taken by the GSC
one-day operation on 2016 November 1. The color gradients represent the event rate, mostly background, perunit area. The most of the entire sky is covered, butthere are some extraordinary areas in the left side. The
thick area is covered by the GSC 3 counter whose anticoincidence background cut is malfunctional, and thus
the rate is high. The thin area is by the GSC 0 counter
with a gas leak.

(South Atlantic Anomaly) and near the geomagnetic
poles. Considering the extension operation, we hold theredundant counter at the stand-by mode. Three gas
counters in GSC 3, 6, and 9 units are considered to havea broken anode wire. They are operated under the limited functionality. The counter in GSC 0 has a gas leak.

2.2. Operation duty cycle
Figure 2 top panels show the daily operation duty cycles of 12 gas counters, GSC 0, 1, 2, ..., B. Because
the carbon-anode wires are consumed by signal charges,
we reduced the anode voltage from the initial 1650V to
1550V. They are typically ∼ 40% because the operation is stopped in the high radiation area in the SAA

Signal gain in gas counter and electronics are monitoredwith 55 Fe 5.9 keV line. Figure 2 bottom panels show the
relative variation of the pulse heights of 5.9 keV line ineach of 12 counters. Except for the GSC 0, the gain ofworking counters are stable within 1 % for the 7 years.
The gain of the GSC 0 suddenly began to increase on
2013 June 10. It increased by 60 % for a month.

Fig. 1. GSC sky map taken for a day on 2016 November 1.

2.3. Gas gain stability

32

2011

2012

2013

2014

2015

2016

GSC_0
1550V

1

2010

2011

2012

2013

2014

2015

2016

GSC_6

0.8
0.6

0.4

0.4

0.2

0.2

1

GSC_1

0.8

1
0.6

0.4

0.4

0.2

GSC_7

0.8

0.6

HV (V)

2010

1650V

0.2

1

GSC_2

0.8

1
0.6

0.4

0.4

0.2

0.2

1

1

GSC_3

0.8

GSC_9

0.8

0.6

0.6

0.4

0.4

0.2

GSC_8

0.8

0.6

0.2

1

GSC_4

0.8

1
0.6

0.4

0.4

0.2

GSC_A

0.8

0.6

0.2

1

GSC_5

0.8

1
0.6

0.4

0.4

0.2

0.2
55500

56000

56500

57000

GSC_B

0.8

0.6

57500

Gas P (atm)

0.6

55500

56000

56500

57000

57500

MJD

55

MJD

Fe 5.9 keV PHA (ch)

1
0.8

2011

2012

2013

2014

2015

2016

2010

2011

2012

2013

2014

2015

GSC_0

1.02
1.01
1
0.99
0.98

GSC_6

GSC_1

1.02
1.01
1
0.99
0.98

GSC_7

1.02
1.01
1
0.99
0.98

GSC_2

1.02
1.01
1
0.99
0.98

GSC_8

1.02
1.01
1
0.99
0.98

GSC_3

1.02
1.01
1
0.99
0.98

GSC_9

1.02
1.01
1
0.99
0.98

GSC_4

1.02
1.01
1
0.99
0.98

GSC_A

1.02
1.01
1
0.99
0.98

GSC_5

1.02
1.01
1
0.99
0.98

GSC_B

55500

56000

56500

MJD

57000

57500

1.6
1.4
1.2
1
0.8
0.6
0.4
0.2
0
6000
5000
4000
3000
2000
1000

2016

1.6
1.5
1.4
1.3
1.2
1.1
1
1.02
1.01
1
0.99
0.98

Residual PHA (ch)

2010

1700
1600
1500
1400
1300
1200
1100
1000
900

400
200
0
−200
−400
56400

55500

56000

56500

57000

57500

MJD

56600

56800

57000
MJD

57200

57400

57600

Fig. 3. History of GSC 0 counter anode voltage (top panel), gaspressure in the best-ﬁt viscous ﬂow model (2nd), pulse heightfor 55 Fe 5.9 keV X-rays and that expected from the gas pressureand gain models (3rd), data-to-model residuals for the 5.9 keV(bottom) from 2013 April 18 (MJD=56400) to 2016 September
19 (MJD=57650).

3. Calibration status
We calibrate the energy response function using spectra
for the Crab nebula. As an example, ﬁgure 4 shows the
Crab nebula spectra by the GSC 0 counter every 200
days since the gas leak started, and the ratios to that
expected by the normal counter. We conﬁrmed that theobserved eﬃciency decrease agrees with that estimatedfrom the gas gain increase (section 2.4.) within the error.

10-1

MJD=56000-56200
MJD=56200-56400

10-2

MJD=56400-56600
MJD=56600-56800
MJD=56800-57000-3

MJD=57000-57200

10

MJD=57200-57400

1.0

Ratio

2.4. Gas pressure change in GSC 0
Figure 3 shows the anode voltage, gas pressure model,
pulse height for 55 Fe 5.9 keV line, and the residuals. We
tried two gas-leak model, a free molecular ﬂow represented by exponential decay and a viscous ﬂow by inverse proportional. We employ the Diethon’s gas-gain
model (Knoll 1989), and then ﬁt to the data to determined the model parameters. The results indicate thatthe viscous-ﬂow model better represents the data. The
data-to-model discrepancy get larger after one and halfyear later. This may represent some physical process
may change. The model suggest that the gas pressure is
now at about 0.1 atm.

Normalized counts s-1 keV-1

Fig. 2. GSC operation duty cycle (top) and gas-ampliﬁer gain stability(bottom). Blue and red lines represent the operation anode voltages of 1650 V and 1550, respectively. Gray terms are operated
in the limited functionality.

0.5

0.0

2

5

10

20

30

Energy (keV)

Fig. 4. GSC 0 Crab nebula spectra for every 200-d interval from2012 March 14 (MJD=56000) to 2016 January 22 (MJD=57400)(top) and their ratios to the model (bottom), calculated fromthe standard Crab-nebula spectrum and the response function of
normal GSC counters with 1.4-atm Xe gas.

References
Knoll, G. K. 1989, Radiation Detection and Measurement
Mihara, T., et al. 2011, PASJ, 63, 623
Sugizaki, M. et al. 2011, PASJ, 63, 635

33

The Current Status of Solid-state Slit Camera aboard MAXI
Hiroshi Tomida,1 Hiroshi Tsunemi,2 Satoshi Nakahira,3 and MAXI/SSC team1

Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency,
3-1-1 Yoshinodai, Chuo, Sagamihara, kanagawa, 252-5210 Japan2
Department of Earth and Space Science, Osaka University,
1-1 Machikaneyama, Toyonaka, Osaka 560-0043 Japan
3
Human Spaceﬂight Technology Directorate, Japan Aerospace Exploration Agency,
2-1-1 Sengen, Tsukuba, Ibaraki, 305-805 Japan
E-mail(HT): tomida.hiroshi@jaxa.jp
Abstract

We present the current status of the performance of the Solid-state Slit Camera (X-ray CCD camera of
the MAXI mission). The charge transfer eﬃciency (CTE) of the CCDs have been degrading, and the energy
resolution has been also degrading. On the other hand, the cooling capability of the MAXI/SSC has beenstable since the beginning of the in-orbit operation. The radiator temperature and camera temperatureshows no signiﬁcant change without those coming from the Sun-orbit plane angle, and the cooling powerof the Peltier devices has been unchanged. Then, the CCD temperatures have been stable.
Key words: MAXI: SSC — X-ray Camera: CCD

1. CCD performanceFig.1 represents the spectrum of Cu-kα line obtainedwith SSC-H/CCDID-0, that are those applied for withthe correction of CTE (charge transfer eﬃciency) andtemperature-gain dependency. Fig.2 represents the longterm trend of the pulse height of Cu-kα line and the line
width of all CCDs. We can see gradual broadening of theline width and the shift of the Cu-line center. The causesof the width broadening are degradation of the chargetransfer eﬃciency, and inaccurate gain-corrections. The
performance degradation is more serious for the CCDs
whose position is close to the camera slit. The intensityof charged particles is higher at the position closer tothe slit, which likely causes the deference of the degradation level. The energy resolution of SSC-Z/CCDID-0
is very bad, which is due to the insuﬃcient cooling powerof the Peltier cooler. Then, we should not use the dataof SSC-Z/CCDID-0. In the MAXI operation, the dataof SSC-Z/CCDID=0 have not been down-linked. Between Mar 2013 and Apr 2014, SSC-Z was not operated
due to the on-board-computer trouble. After the nonoperation period, SSC-Z has functioned completely. The
performance of the SSC-H and Z had been gradually degraded. For the detailed spectrum analysis of SSC, weshould be careful when you use the data obtained after2014.

Fig. 1. Cu-Kα line proﬁles of SSC-H/CCDID-0. 6 spectra are extracted from single-pixel event data in 6 diﬀerent periods, thatare shown above each panel.

34

Fig. 3. Temperature trend of the radiator panels (left) and SSC
camera body (right). Each dot represents the one-day averaged
temperature. MAXI has two radiator panels.

Fig. 2. Trend of characteristics of Cu-kα line. The center and width of
the Gaussians are plotted. Each dot (spectrum) is extracted fromthe 3-months integrated data. In each panel, data of 16 CCDs
are plotted.

2. Cooling Capability
CCD chips of SSC are cooled with Peltier coolers in the
CCD devices, and the heat from the Peltier coolers aretransported to radiator panels on the surface of MAXI
payload via the loop heat pipe. Fig.3 (left) shows thelong term trends of the radiator temperature, and Fig.3(right) shows that of SSC camera body. The orbital
plane of the International Space Station precess every 70
days, then the temperatures show the periodic change according to the precession. However no inclination can be
seen. There is no degradation for the cooling capabilityof the radiator and loop heat pipe system in the 7-years
MAXI in-orbit operation. The temperature change ofthe camera body around 2014 corresponds to the periodthat SSC-Z was not operated due to the computer trouble. During this period, the heat generation from SSC
was small compared with that in other period. Fig.4
shows the temperature diﬀerence of CCD chips and camera body, which represents the colling power of Peltier
devices. Peltier Current is 1.0 A and 0.82 A for SSC-H
and SSC-Z, respectively. No signiﬁcant degradation canbe seen. A small change around 2014 in SSC-H is due tothe temperature increase pf camera temperature when
SSC-Z was not operated.

Fig. 4. Cooling power of 32 Peltier devices. The vertical axis
shows the temperature diﬀerence between CCD ship and camera
body. The left upper, right upper, left bottom, and right bottom panel are the data of SSC-H/CCDID0-7, SSC-H/CCD8-F,
SSC-Z/CCD0-7, and SSC-Z/CCD8-F, respectively.

35

Energy spectra of bright clusters of galaxies
obtained from the MAXI data archive
Tae Furusho,1 Yasuharu Sugawara, 1 Satoshi Nakahira,3 and Ken Ebisawa1

1

Japan Aerospace Exploration Agency, Institute of Space and Astronautical Science,
3-1-1 Yoshino-dai, Chuo-ku, Sagamihara, Kanagawa 252-5210, Japan2
JEM Mission Operations and Integration Center, Human Spaceﬂight Technology Directorate,
Japan Aerospace Exploration Agency, 2-1-1 Sengen, Tsukuba, Ibaraki 305-8505, Japan
E-mail(TF): furushoh.tae@jaxa.jp
Abstract
In order to enhance use of the MAXI data by scientists all over the world, we are developing the MAXI
data archive. The archival data and tools will be released from DARTS at ISAS/JAXA by early 2017.
Also, it is planned that the data and tools be available from HEASARC at NASA/GSFC. Although MAXI
data analysis procedure is complicated, we oﬀer simple analysis tools that enable end-users to easily create
standard images, energy spectra and responses, and light curves.
Key words: workshop: MAXI data archive — Clusters of galaxies

1. MAXI data archive
Long-time MAXI observations are scientiﬁcally valuable
data for astrophysics. MAXI Archive data are planned
to be released from DARTS at ISAS/JAXA in Feb 2017,
and from HEASARC at NASA/GSFC in 2017. Becauseof the peculiar system of MAXI, we oﬀer simple analysis
tools for scientists, even who are not familiar with thedetector details.
1.1. Design of MAXI data archive system
The MAXI archive data are processed and screened using the latest CALDB and auxiliary data. Users do notneed screening. When new CALDB is released, archive
data are updated. Daily event ﬁles are divided into
768 Healpix regions. Users can download event ﬁlesand auxiliary ﬁles corresponding to a speciﬁed sky region. Or, users can locally have the entire MAXI archive( 800 Gbytes for seven years). Data archive are updated
real-time (typical delay 15 minutes), and released fromDARTS at ISAS/JAXA, and mirrored to HEASARC at
NASA/GSFC. Using MAXI speciﬁc ftools and CALDB(will be included into HEAsoft), users can analyze data
easily. Figure 1. shows the structure of MAXI data
Archive. Event data are put into a directory of eachday, MJDxxxxxx/.
1.2. MAXI ftools - mxftools
MAXI analysis ftools consist of nine MAXI ftools, andone script for data download. Users basically use only

mxdownload and mxproduct. mxdownload downloads data from the archive for speciﬁed coordinates.
mxproduct is a perl script running the MAXI ftools
to generate images, light curves, spectra, and response
ﬁles, calling the following tools.
mxextract Create a single standard event ﬁle from plural region event ﬁles for a given coordinates (RA,DEC) and radius.
mxgscancur, mxsscancur Create an observation andinstrument speciﬁc ﬁle to describe observational
conditions.
mxgscandat, mxsscandat Create observation and instrument speciﬁc scan history ﬁles for GSC and SSC.
mxgtiwmap Create an exposure histogram for each incident angle of a given input target.
mxlcscan Create light curve ﬁles.
mxrmfgen Create a response ﬁle for GSC and SSC.
1.3. How to use analysis tools
First, you must have HEAsoft, CALDB, and also Perl,
Python, Ruby. And then, install MAXI software package, mxftools, and set it up. Second, download events/
and auxil/ data ﬁles for your target using a command,
mxdownload, with parameters of RA,DEC, start stop
time, and radius:% mxdownload RA DEC tstart tstop radius--uri=https://maxi.darts.isas.jaxa.jp/mxdata

36

obs/
|
MJD55000/ MJD56000/ MJD57000/ (every 10000 days)|
MJD55000/ MJD55001/ MJD55999/ (every 1 day)|events/ auxil/||_( files for attitude, orbit, time, ISS data, GSC time data, MKF, HK etc.)
gsc_med/ gsc_low/ ssc_med/( for GSC and SSC, for bit rate med/low )|||_ mx_mjd55000_ssch_med_000.evt ... mx_mjd55000_sscz_med_767.evt||_ mx_mjd55000_gsc_low_000.evt ... mx_mjd55000_gsc_low_767.evt
mx_mjd55000_gsc_med_000.evt ... mx_mjd55000_gsc_med_767.evt( event files for 768 Healpix regions)Fig. 1. Directory and ﬁle structure of MAXI archive.

30:00.0

25:00.0

20:00.0

15:00.0

05:00.0

10:00.0

3:00:041:00:00.0

42:00:00.0

43:00:00.0

35:00.0

normalized counts sí1 keVí1

0.02

5×10í3

2×10í310í3(dataímodel)/error

40:00:00.0

4

39:00:00.0
42

0.01

1066

1190

1316

1440

1566

1690

1814

1940

2064

21

2
0
í2
5
Energy (keV)

10Fig. 2. Output GSC image and spectrum of Perseus cluster for 6 years.

If users already have local mount data, this step canbe skipped. Third, run mxproduct with 4 parametersof RA, DEC, start, stop:% mxproduct [options] RA DEC tstart tstop
You can set various options, such as object name, energy bands for light curves, regions, GSC or SSC only,
etc. Output ﬁles, event ﬁles, images (all, wmap), energy spectra (source, background), responses, and light
curves are created. Also, you can leave intermediate ﬁles
if you want to know detailed analysis.
2. Examples of analysis
Here we show an analysis example of the Perseus cluster. Figure 2 shows the GSC image and spectrum of
the Perseus cluster for 6 years, 2011–2016. The best-ﬁt
parameters agree with the MAXI spectrum by MAXI on-

demand analysis (http://maxi.riken.jp/mxondem/),
as well as with those by other missions. Thanks to
MAXI’s sensitivity to largely extended emission andwide energy band, we are able to create high-quality
spectra in 0.7–20 keV from the entire clusters of galaxies.

37

MAXI Transients

38

39

The Swift and MAXI Galactic Transient Collaboration
Jamie Kennea,1 Hitoshi Negoro,2 Phil Evans,3 Andrew Beardmore,3Hans Krimm,4 Motoko Serino,5 Patrizia Romano,6 and Kazutaka Yamaoka7
1

Dept. of Physics and Astrophysics, The Pennsylvania State University, University Park, PA 16802, USA
2
Dept. of Physics, Nihon University, 1-8-14 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8308, Japan
3
Dept. of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK
4
National Science Foundation, 4201 Wilson Blvd, Arlington, VA 22230 USA
5
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
6
INAF/IASF Palermo, via Ugo La Malfa 153, I-90146 Palermo, Italy
7
Dept. of Particle Physics and Astronomy, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8601, Japan
E-mail(JK): jak51@psu.edu
Abstract
NASA’s Swift and JAXA’s MAXI X-ray telescopes have, for the past 7 years, collaborated on the scienceand discovery of X-ray Binary Transients in our Galaxy and in the nearby Small and Large Magellanic
Clouds. MAXI’s near all-sky coverage every 90 minute ISS orbit, in the 0.5-20 keV energy range, allows
for the rapid detection of bright (> 40 mCrab) X-ray transients, and through the MAXI Nova Alertsystem, rapidly reports these new transients to the community. Swift’s capability of rapid slewing andautonomous TOO follow-up, combined with the well matched ﬁeld of view of the Swift/XRT allows for
accurate localization and follow-up of MAXI discovered transient sources. Here we discuss in detail how
MAXI and Swift work together to localize X-ray transients, and the detail recent science results on SMC
X-3 and the short lived transient sources MAXI J1957+032 and MAXI J0636+146.
Key words: workshop: 7 Years of MAXI

1.

Introduction

The “Monitor of All-sky X-ray Image” (MAXI; Matsuoka et al. 2009) is a near all-sky imaging X-ray telescope installed on the Japanese Experiment Module Exposed Facility (JEM-EF) on board the International
Space Station (ISS). It consists of two X-ray detectors,
the the Gas Slit Camera (GSC) and the Solid-state Slit
Camera (SSC). The GSC consists 6 one dimensional position sensitive proportional camera, and provides sensitivity in the 2-30 keV energy range, with a FOV of 1.5
deg x 160 deg. The SSC is Charge Couple Device (CCD)
based X-ray detector, operating in the energy range 0.5- 12 keV, consisting of 32 CCDs with a 1.5 deg x 90 deg
FOV. These cameras are ﬁxed with respect to the ISS,and scan the sky as the ISS completes its 90 minute orbit, building up an almost complete image of the X-ray
sky.
MAXI s ability to collect an X-ray image of the sky
in the 0.5-20 keV energy band, with sensitivities as low
as 60 mCrab (5 sigma) in a single orbit and 15 mCrab
in a day, makes it extremely capable at ﬁnding bright
new transients in the local universe, in particular in thispaper, we focus on X-ray binary sources discovered to bein outburst by MAXI in our Galaxy and in the nearby
Small Magellanic Cloud (SMC).
The MAXI mission began when it was launched onboard the Space Shuttle and mounted on JEM-EF onboard the ISS on July 24th, 2009, and has been operating
continuously since then.
NASA’s Swift mission (Gehrels et al. 2004) was
launched in November 2004 with the primary science
goal of detecting and localizing Gamma-Ray Bursts(GRBs).
Swift consists of three instruments, thehard X-ray (15-150 keV) Burst Alert Telescope (BAT;
Barthelmy et al. 2005), a coded mask large area monitor
capable of imaging 1/6th of the sky, the X-ray Telescope(XRT; Burrows et al. 2005), a Wolter Type-I X-ray (0.3
- 10 keV) telescope with CCD camera, with a ∼ 24 arcminute diameter FOV, and the Ultraviolet/Optical Telescope (UVOT; Roming et al. 2005), a Ritchey-Chrétien
optical telescope based upon the XMM-Newton Optical
Monitor, with a 17 arc-minute square FOV, 6 ﬁlter (3optical, 3 UV) imaging, and optical and UV grisms forspectroscopy.
The combination of these three instruments, alongwith a robotic spacecraft platform that allows for very40 

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Fig. 1. Example of XRT localization of MAXI J0556−332 with Swift.
MAXI error circle is shown over XRT Photon Counting mode data,clearly showing a bright transient source inside the MAXI error
circle.

rapid slewing, and automated follow-up of on-board detected GRBs as well as and ground command Target-ofOpportunity (TOO) observations make Swift the ideal
mission not only for detecting GRBs, but also for detecting, localizing and following up a wide range of transientphenomenon.
For X-ray binaries, Swift/BAT is able to localize anddetect the rising brightness of hard X-ray sources, making it especially capable of detecting black hole candidate(BHC) low mass X-ray Binaries (LMXBs), but less capable of detecting LMXBs containing neutron star (NS)compact objects, as these often have spectra that cut oﬀ
above 10 keV. MAXI’s lower energy response makes itmuch more sensitive to these transients, and therefore
acts as a great compliment to the capabilities of BAT.
In this paper we discuss how utilizing Swift is utilized
to localize and provide long-term monitoring of MAXItransients, and discuss recent science results as a resultof this collaboration.
2. Swift response to MAXI transients2.1. Target of Opportunity observations
NASA’s Swift observatory has several unique featuresthat make it an ideal observatory for following X-ray
transients. For transients discovered on-board by the
BAT instrument, Swift will automatically slew to thearc-minute accuracy BAT position, and perform automated observations of the source, reporting rapidly via
the Tracking and Data Relay Satellite System (TDRSS)
the position, spectrum and light-curve of the object, in

hard X-ray, X-ray and optical/UV. XRT can localize apoint source to an accuracy of 3.5 arc-second radius (alllocalization accuracies in this paper quoted at 90% conﬁdence). Furthermore utilizing UVOT images of the ﬁeld
to correct for the systematic errors in the star tracker’s
pointing solution, this error can be reduced to as good as
1.5 arc-second radius (90% conﬁdence). In many cases
where transients have low absorption, UVOT may also
detect an optical/UV counterpart to the X-ray transient,
allowing for a sub-arc-second localization of the source,
and a broad-band spectral energy distribution (SED) tobe created utilizing data from optical, UV, X-ray andhard X-ray.
Swift is a robotic telescope and has an automated
TOO feature that allows for ground commanded TOO
observations, that override observations in the daily preplanned science timeline. These TOO commands can be
performed rapidly, either through TDRSS (with a 30
minute latency) or through ground station passes that
occur regularly throughout the day. Results of such
TOO observations are typically available within hours ofobservation, depending on the timing of ground station
contacts after the upload. Streamlining and automation
work by the Swift Science Operations Team means thatsetting up Swift TOO observations is a simple task, utilizing a web interface for submission of TOO requests,
and a simple question and answer based command line
tools for setting up TOO uploads, than can be performed
over any secure Internet connection.
Swift TOO requests can be requested a variety of priorities: Priority 1 (“Highest Urgency”) TOOs require
rapid follow-up within 4 hours if possible, and will ifnecessary wake up the on call observatory duty scientist(ODS) to alert them that an observation needs to be
performed; Priority 2 (“High Urgency”) requires observations within a 24 hour period, and pages the ODS if it
is within working hours (0800 - 1700 Eastern time); Priority 3 (“Medium Urgency”) if observations are needed
within days to a week; Priority 4 (“Low Urgency”) for
longer term observations (weeks to months).
Swift TOO requests are open to the entire community
and in make up a large fraction of the observations that
Swift performs. In 2016 there were 1300 approved TOO
requests for Swift, at varying levels of priorities, and inmany ways Swift has evolved to become an observatory
where observations of transient phenomena of all kinds
through TOOs has become the primary science.
2.2. Swift’s response to MAXI transients
MAXI’s GSC, high sensitivity and low background is themain instrument that is utilized for detection of X-ray
transients. SSC suﬀers from background issues, making it diﬃcult to use for transient searches. X-ray transients in MAXI data are detected by the “Nova-Alert
System”, as detailed by (Negoro et al. 2016). Detec-

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Fig. 2. Swift/XRT X-ray light-curve of SMC X-3 during its outburst
in 2016-2017.

tions are disseminated via email alerts to the X-ray transient community, and later via Astronomers Telegram(ATEL). The typical error region for a slowly varying
X-ray transient detected by MAXI is ∼ 0.2 degrees radius, although fainter transients, or transients that areundergoing rapid variability on the timescale of a MAXItransit will have larger error regions. This 0.2 deg error
region is well matched with the Swift/XRT FOV, whichis also approximately 0.2 degrees radius, meaning thatmost bright MAXI transients can be localized with a single pointing of the Swift XRT. An example of this typeof localization is given in Figure 1.
For fainter transients, where the source is varying over
the transient time of MAXI, or in crowded regions of thesky (e.g. the Galactic Center region) the error regions
are typically larger than can be covered by a single XRTpointing. For these cases Swift is able to perform automated tiling observations of approximately circular regions. These tiling observations consist of 4, 7, 19 or 37
tiles. Although the larger number of observations meansthat each tile must be exposed for a shorter exposure
time, the brightness of MAXI detected transients (> 5
mCrab) equates to XRT count rates of > 3 count/s,
which given the very low background in the XRT, areeasily detectable in observations as short as 60s.
The Swift/MAXI transient group have had a program,approved through the Swift Guest Investigator program,for high priority rapid follow-up of MAXI detection Xray transient sources in the Milky Way and nearby galaxies (i.e. the SMC and LMC). This program has been
approved for Swift Cycles 6, 7, 8, 9, 11, 12 and 13, covering the period of April 1st, 2010 to March 31st, 2018.

Although the program was not approved for Swift Cycle
10 (April 1st, 2014 to March 31st 2015), observations
during this time were performed as part of the regular
TOO program.
This program has the goal of providing localizations
and identiﬁcations of MAXI transients, and the program
team comprised both members of the Swift and MAXI
teams. As part of this program, the Swift/MAXI transient team has submitted 55 TOO requests for MAXI
follow-ups, and reported results in at time of writing, 56
ATELs. In addition, many papers have been publishedby the Swift/MAXI transient group and others as a result of the discoveries made by MAXI and this program,and utilizing the follow-up observations triggered via the
Swift TOO program, including the discoveries of 6 new
BH candidates by MAXI: MAXI J1659−152 (Kennea etal. 2011), MAXI J1836-194 (Ferrigno et al. 2012), MAXI
J1543−564 (Stiele et al. 2012), MAXI J1828−249 (Filippova et al. 2014), MAXI J1910−057 (Degenaar et al.
2014) and MAXI J1305−704 (Shidatsu et al. 2013).
Beyond this initial localization, the Swift TOO program has been utilized in order to perform monitoring
observations of many of the MAXI sources.
3. Recent Science Results
In this section we will highlight results of science resultsthat have been enabled by the Swift/MAXI collaboration. Due to the large number of results as a result of
this program it would be impossible to list all the science
results from this collaboration during the entire MAXI
mission in this paper, so instead we highlight a few results obtained in the last 12 months of this collaboration.
We also do this in order to highlight that the collaboration between Swift and MAXI remains highly productive
one, even given the relative old ages of the two missions.
3.1. SMC X-3
SMC X-3 is a Be/X-ray binary in the SMC, with an optically measured orbital period of 44.86 day (Cowley andSchmidtke 2004). In addition it has a measured pulsar
period of 7.78s (Edge et al. 2004a). It has demonstrated
large Type II outbursts (outbursts not linked to the orbital period) on several occasions since it was ﬁrst discovered by the SAS 3 X-ray observatory in 1977 (Clark
et al. 1978).
On 13:17UT on August 8th, 2016, MAXI detected abright X-ray transient in the SMC. As the SMC is adense source of X-ray transients, several of which werein outburst at the time, it was not possible to positively
identify at the time which source was in outburst, so
this outburst was named MAXI J0058-721 (Negoro etal. 2016b).
Observations performed on August 10th, 2017 by Swift
XRT found a bright point source at a position consis-

42

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Fig. 3. Pulsar period of SMC X-3 over time, with a model for both spin-down based on the equations of Ghosh and Lamb (1979), and orbital
modulation due to the Doppler eﬀect.

tent with SMC X-3, and inside the error circle for MAXI
J0058−721 (Kennea et al. 2016a). The XRT position
is consistent with the Chandra position of SMC X-3 reported by (Edge et al. 2004b). Given the appearance of
SMC X-3 in such a bright state, the brightest it has been
seen since its previous outburst in 2003, it was concluded
that that MAXI J0058−721 was indeed the reappearanceof SMC X-3.
Triggered by the MAXI detection of SMC X-3, Swiftbegan a series of monitoring observations to track theoutburst of SMC X-3. As SMC X-3 was bright, and had aknown pulsar period, XRT data was taken in Windowed
Timing mode (WT) in order to avoid pile-up and to allow us to obtain an accurate measure of the pulsar period, which would provide another conﬁrmation that thissource is indeed SMC X-3. These timing observations did
indeed conﬁrm this, with observations taken on August
10th, 2016 showing a period of 7.812749 ± 0.000012 s(Kennea et al. 2016b). These observations also showedthat the source was hard, well ﬁt by a power-law modelwith a photon index of 1.10 ± 0.02, and when corrected
for a standard SMC distance of 61 kpc, the X-ray luminosity was 3.7 × 1038 erg/s/cm2 (0.5-10 keV), a value
which is > 2LEdd , without correcting the lumiinosity tobolometric. SMC X-3 was undergoing super Eddington
accretion.
Table 1. The orbital parameters of SMC X-3 as derived by ﬁtting
the evolution of the pulsar period throughout the 2016–2017 outburst to a combined model of spin-up model of Ghosh and Lamb(1979) and a model of Doppler shifted orbital modulation of thepulsar period. For details of this orbital determination please see
Townsend et al. (2017)

Parameter
Orbital Period
Eccentricity
Omega
ax (sin)(i)
TperiastronValue
45.383 ± 0.421 days
0.22 ± 0.01
208.77 ± 2.86 degrees
196.2 ± 3.7 light-s
MJD 57676.90 ± 0.56

Over the next few months, Swift monitored both theﬂux and pulsar period evolution of SMC X-3. After ashort period spent rising to a it reached a peak X-ray
luminosity of ∼ 1039 erg/s/cm2 , after which followed aslow decline of ﬂux. At the time of writing SMC X-

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Fig. 4. X-ray light curve of MAXI J1957+032 as seen by MAXI (star)and Swift (circles).

3 is still in outburst, but at brightness approximately
2 orders of magnitude below its peak. The count rate
evolution of SMC X-3 can be seen in Figure 2.
The pulsar period evolution of SMC X-3 (also see Figure 3), appears complex. Overall the pulsar is as expected spinning up during the accretion, but the rateis complex and there appears to be some modulation.
Given that the SMC X-3 has a measured 45 day period, we attempted to ﬁt this light-curve with a model ofspin-up and orbital period induced changes in the pulsar
period due to the Doppler eﬀect. However, this model
provided a poor ﬁt. In an attempt to better model thischange, and given the extreme nature of the accretion
onto SMC X-3, we decided to model the spin-up using
the model of Ghosh and Lamb (1979), following from thework of Takagi et al. (2016), where spin-up (Ṗ ) is modeled based on the X-ray luminosity of the source. The
ﬁtted orbital parameters of SMC X-3 are given in Table
1. As can be seen the orbital period measured from thismethod matches withing errors that measured in onoptical by Cowley and Schmidtke (2004), and in addition,
we for the ﬁrst time have a measure of the orbital eccentricity of SMC X-3. The full results of this modeling are
and further analysis of the 2016–2017 outburst of SMC
X-3 are reported by Townsend et al. (2017).
3.2. MAXI J1957+032
MAXI J1957+032 is a faint X-ray transient ﬁrst detected
on May 11th, 2015 by MAXI/GSC (Negoro et al. 2015).
It was also detected by INTEGRAL (Cherepashchuk et
al. 2015) and seen to decay quickly(Molkov et al. 2015).
Further repeated outbursts occurred on October 6th, 

Fig. 5. Spectral evolution of MAXI J1957+032 with ﬂux, as seen fromall Swift/XRT data taken from the three observed outbursts.

2015 (Sugimoto et al. 2015), January 7th, 2016 (Sugimoto et al. 2015) and on September 29th, 2016 (Negoro
et al. 2016a). Swift performed rapid follow-up observations of three out of four of these triggers.
MAXI and Swift observations show similar outburst
proﬁles for all three observed triggers, that of a rapid
decay within a period of 6 years from detection. We
show in Figure 4 the outburst light curve of the most recently observed outburst oﬀ September 29th, 2016. The
outburst light curve can be well described by a period ofslow decay lasting ∼ 2 days, followed by a rapid decline
to non-detection over the next 3–4 days. The rapid decline is somewhat similar to what is seen from X-ray pulsars rapidly decay due to the propeller eﬀect (e.g. Tsygankov et al. 2016), although other explanations cannot
be ruled out as to the rapid decline. Also there is no evidence of X-ray pulsations is found in MAXI J1957+032
XRT data.
During these outbursts the source also shows considerable spectral variability. Shown in Figure 5, we plot
the ﬁtted photon index for all the Swift observations of
MAXI J1957+032, plotted against ﬂux, and ﬁnd thatthey are anti-correlated, with r = −0.91. This softening of spectrum has been reported to be seen in NS
X-ray binaries by Wijnands et al. (2015), and the evolution of the photon index with ﬂux is very similar to thatshown in Figure 1 of that paper for NS sources. Assuming this is the case, this points to a NS origin for MAXIJ1957+032, and scaling the ﬂuxes to the luminosities
from that ﬁgure, allows us to estimate the distance to
MAXI J1957+032 as being between 2-4 kpc. We plan to44follow future outbursts of this source with longer exposures in order to search for signatures of a pulsar and tobetter determine the shape of the outburst light-curve.
3.3. MAXI J0636+146
On November 4th, 2016 at 02:20UT MAXI detected a
new X-ray transient, named MAXI J0747+146. This
source was seen at a level of 51 ± 13 mCrab for two
orbits, and then was no longer detected, proving it tobe short lived (Negoro et al. 2016c). Swift observations rapidly followed at 06:12UT, less than 4 hours after
the MAXI detection. Swift localized the transient anddetermined that it was not a previously known point
source. UVOT observations showed no detection of theobject (Kennea et al. 2016c). Swift monitoring observations were triggered, until the source was no longer detected by Swift 5 days later. The rapid decline of MAXI
J0636+146 prompts speculation that it may be similarsource type to MAXI J1956+032. However, several important diﬀerences are apparently, ﬁrstly the spectrum of
the source is very soft, best ﬁt with a black-body modelwith kT = 0.6 ± 0.1 keV, furthermore there is not strong
evidence for spectral evolution throughout the outburst,
although time resolved spectral ﬁts are consistent withcooling, they are also consistent with a constant level
temperature, so better quality spectra would be needed
to make any strong conclusion on this. Given the lack of
optical counterpart, especially as the ﬁtted X-ray absorption is essentially negligible, the origin of this outburst
remains a mystery.
4. Conclusion
Swift and MAXI have proven to be a highly successful combination for the rapid detection, localization andfollow-up of X-ray transients in the local Universe. Bothhave unique capabilities, MAXI with its near all-sky coverage in soft X-rays, and Swift with its sensitive focused
instruments and ability to rapidly get on-target, and also
perform long term follow-up observations.
5. Acknowledgements
J. A. Kennea acknowledges support from NASA aspart of the Swift Guest Investigator Program for thiswork. The current work is supported under NASA Grant
NNX16AR15G.
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45

The orbital period of MAXI J1305-704
Aarran W. Shaw,1,2 Philip A. Charles,2 Jorge Casares,3,4,5
Danny Steeghs6
1

Department of Physics, University of Alberta, CCIS 4-181, Edmonton, AB T6G 2E1, Canada
School of Physics and Astronomy, University of Southampton, Southampton SO17 1BJ, UK
3
Instituto de Astrofı́sica de Canarias, E-38205 La Laguna, S/C de Tenerife, Spain
4
Departamento de Astrofıisica, Universidad de La Laguna, E-38206 La Laguna, S/C de Tenerife, Spain
5
Department of Physics, Astrophysics, University of Oxford, Keble Road, Oxford OX1 3RH, UK
6
Department of Physics, Astronomy and Astrophysics Group, University of Warwick, CV4 7AL Coventry, UK
E-mail: aarran@ualberta.ca
2

Abstract
MAXI J1305-704 is a candidate black hole X-ray binary discovered by MAXI in 2012. The systemparameters are relatively unknown, with no dynamical determination of the mass of the compact object
as of yet. There have also been several conﬂicting measurements of the source’s orbital period, the mostrecent of which is Porb = 9.74h, based on the observation of deep X-ray dips in the X-ray light curve. We
reanalyze a 40 ks Suzaku observation of MAXI J1305-704 in order to disentangle the multiple periodicities
previously reported. Using timing analysis techniques, we suggest that the orbital period is instead ∼ 5h.
We also present optical spectroscopy of the source in outburst that provides a mass function, but its value
is too low to constrain the nature of the compact object.
Key words: black hole physics — X-rays: binaries — X-rays: individual: MAXI J1305-704

1.

Introduction

Galactic black hole X-ray transients (BHXRTs) are lowmass X-ray binaries consisting of a black hole (BH) accreting from a low-mass (≤ 1M ) companion. They
are characterized by long periods of quiescence (years to
decades) followed by X-ray outbursts that can increase
the luminosity by several orders of magnitude. BHXRTs
have proven to be important in studying X-ray binaries(XRBs) as in quiescence they provide the opportunity
to detect the donor itself, which is mostly impossible inluminous, persistent XRBs (Charles & Coe 2006)
Determining the system parameters of BHXRTs is
crucial for understanding their formation and evolution.
However, it is challenging as most systems are too faintin quiescence for radial velocity studies using the current
generation of large telescopes. A key parameter requiredfor determining the mass function (and therefore the nature) of the compact object is the orbital period (Porb )of the system, which can often be measured in outburst(see e.g. Zurita et al. 2008).
Since its launch in 2009 MAXI /GSC (Matsuoka et
al. 2009) aboard the ISS has discovered 6 new candidate BH systems. One such system, MAXI J1305-704(hereafter J1305), was discovered in April 2012 (Sato etal. 2012). Optical follow up revealed a new source notpresent in archival Digitized Sky Survey (DSS) images(Greiner, Rau & Schady 2012) and a spectrum typical of
an LMXB in outburst (Charles et al. 2012). Pointed observations of the source with Swift revealed a soft, diskdominated spectrum (Kennea et al. 2012a).
The Porb of J1305 is the subject of much debate, owing to the wide range of values that have been reported.
Swift observations of the source during its initial outburst revealed dips in the X-ray light curves accompaniedby spectral hardening. The dips occurred at irregulartimes and made determining the exact periodicity diﬃcult but Porb has been suggested to be 1.5h or 2.7h (Kennea et al. 2012b). Subsequently, Shidatsu et al. (2013)
reported Porb = 9.74 ± 0.04h based on the discovery of
deep dips in Suzaku light curves. However, this value
was determined by simply measuring the time between
the dips, rather than a formal timing analysis. The discovery of dips in multiple X-ray observations indicatesthat J1305 has a high inclination angle, a hypothesis
supported by the presence of strong Fe L absorption features around 1keV in the Chandra/HETGS spectrum of
the source (Miller et al. 2012). However, detailed constraints on the inclination have not yet been obtained.
Here, we summarize our reanalysis of the Suzaku observations of J1305 presented by Shidatsu et al. (2013),

46

focusing in particular on timing analysis of the light
curves. We also present optical spectroscopy of thesource at the peak of its 2012 outburst and study theemission features, attempting to place a limit on themass of the compact object.
2. Observations and Data Reduction
2.1. X-ray observations
We used archival Suzaku/XIS (0.2 – 12keV) observations of J1305 from 2012 July 20 giving ∼40ks (OBSID: 907001010). The observation took place whilst thesource was in a low/hard accretion state (Shidatsu etal. 2013), displaying an average count rate of ≈ 5 cts
s−1 so pileup eﬀects are negligible and can be ignored.
We reduced the data with heasoft v6.19, emulating theanalysis performed by Shidatsu et al. (2013) and using
xselect to generate light-curves in the standard manner.
2.2. Optical observations
J1305 was observed on the nights of 2012 April 15–
16 with the Robert Stobie Spectrograph (RSS) on theSouthern African Large Telescope (SALT; Buckley et al.
2006) at the SAAO. A total of 2.25h of continuous spectroscopy was obtained on each night using the G2300
VPH grating with individual integration times of 600s.
We used a slit width of 0.6 , covering a wavelength rangeof 4040 – 5110 Å.
We used standard data reduction techniques for SALT
as described by Crawford et al. (2010). Wavelength calibration was performed with iraf (Tody 1986) to achieve
the solution for the Xe and CuAr comparison arcs. We
then transformed the 2D spectra on to the wavelengthscale and extracted the 1D spectra with the iraf task
apall.
3. Results
3.1. X-ray Light Curves
The 128s binned Suzaku light curves are presented inFig. 1. Only the XIS-3 light curves are plotted for clarity, but the XIS-0/1 light curves are almost identical.
They exhibit signiﬁcant variability in both bands, withdeep dips appearing at irregular intervals accompanied
by notable increases in the hardness ratio.
It is diﬃcult to estimate the periodicity of the dipsby eye due to data gaps, mostly due to the low-Earthorbit of the spacecraft. We therefore undertook a LombScargle analysis on both the light curves and the HR todetermine a central peak frequency followed by a MonteCarlo simulation based on ‘bootstrap-with-replacement.’
The distribution of peak frequencies resulting from 10000
iterations of this process indicates the error on the peakfrequency. The periodograms exhibit a number of periodicities (Fig. 2), making it diﬃcult to determine thetrue Porb . The strongest peaks correspond to periods of
3.14±0.06h (3.23±0.05h) and 4.98±0.05h (4.99±0.04h)
in the 0.7 – 2keV (2 – 10keV) energy band.
Fig. 3 shows the folded Suzaku light curves on thetwo most signiﬁcant peaks from Fig. 2. The folded light
curves show diﬀering morphology and are explored further in Section 4.1.
3.2. Optical spectrum
The averaged optical spectrum of J1305 (Fig. 4) exhibits strong double-peaked Balmer emission (Hβ , Hγand Hδ ) and He ii λ4686, indicative of a high inclination
binary. Also evident is strong emission in the Bowen region (λ4640 − 4660), a blend of narrow, high-excitation
N iii and C iii emission lines arising from the irradiatedsurface of the donor star.
Closer inspection of the Bowen region reveals that thelines are extremely broad and blended together, making
any estimation of emission line properties (e.g. FWHM)
highly uncertain and is therefore not attempted. The
lack of sharp features in the region makes the Bowen ﬂuorescence technique for measuring the mass of the compact object (Steeghs & Casares 2002) extremely diﬃcult,and is beyond the scope of this work.

4. Discussion and Conclusions
4.1. The Orbital Period
To investigate the multiple periodicities present in the
periodograms in Fig. 2 we must study the light curve
folded on these periods. The phase-folded X-ray light
curves in Fig. 3 display diﬀering morphology dependent
on the period used in the folding. The dipping structure seen in Fig. 1 is recovered in both bands of the∼ 5h folded light curves. Deep dips are seen in the
0.7 – 2keV band, whilst shallower dips appear in the2 – 10keV folded light curve, echoing the apparent dichotomy in dipping structure in the two energy bands as
seen in Fig. 1.
When folded on the other major peak in Fig. 2(3.14/3.23h), the dipping structure is not at all clear in
the 0.7 – 2keV band, though it is visible in the 2 – 10keV
energy band, albeit with large uncertainties. There arealso empty bins in the folded data on these periods. It
is possible that the ∼ 3.2h period can be attributed to
double the spacecraft’s Porb (1.6h; Mitsuda et al. 2007).
There are no peaks in the periodogram at 1.6h, but this
could be due to an irregular observing pattern of thesource over the total ∼ 30.5h observation. The fact thatthe dipping structure is only revealed in both bands ofthe 4.98/4.99h folds of the Suzaku light curves implies
that out of the proposed periods, this period appears tobe the best candidate for the Porb of J1305.

47

0.7-2keV

Flux (count s−1 )

4
3
2
1
0

2-10keV

Flux (count s−1 )3
2
1Hardness Ratio

0

101

100
56128.6

56128.856129.056129.2

56129.4
MJD

56129.6

56129.8

56130.056130.2

Fig. 1. Suzaku XIS-3 light curves in 128s time bins. Plotted are the 0.7 – 2keV (S; Top) and 2 – 10keV (H; Centre) energy bands and thehardness ratio, deﬁned as H/S (Bottom).

50
0.7-2keV

Power

40
30
20
99.99%
10050
2-10keV

Power

40
30
20
99.99%
10
0

0

5

10

15

20

25

30

35

40

45

Frequency (day−1 )Fig. 2. Lomb-Scargle periodograms of the Suzaku light curves in
the 0.7-2keV (Top) and 2-10keV (Bottom). The two most signiﬁcant periodicities seen are marked in each panel with verticaldashed lines. The horizontal dashed lines in each panel represent
the 99.99% signiﬁcance level, determined by randomizing the ﬂux
positions in the light curves whilst maintaining the time stamps
and performing 10000 iterations of the Lomb-Scargle algorithm.

Fig. 3. Suzaku light curves of MAXI J1305-704, folded on the twomost signiﬁcant peaks (from Fig. 2). The top two panels arefolds of the 0.7 – 2keV light curves and the bottom two panels
are folds of the 2 – 10keV light curves. Open points represent data
which have no estimated uncertainties due to only one unbinned
point being available for binning. Also shown in light grey are
the unbinned phase folded light curves in each energy band. The
periods used for folding are designated in each panel.

4.2. System Parameters
The previous estimate of Porb = 9.74h would requirean evolved donor star with M2 , R2 and ρ2 less thanthat of the Sun (Shidatsu et al. 2013). With our revised
Porb = 4.98h we calculate the radius of the donor to
be R2 /R = 0.55 − 0.7 using a sensible range of mass

ratios (M2 /M1 = 0.1 − 0.2; where M2 , M1 are the massof the donor and compact object, respectively) withthe standard approximations for Roche lobe geometry(Eggleton 1983). We also calculate the density ofthe donor to be log(ρ2 /ρ ) ≈ 0.51. These values are

48

Hβ
Bowen

20
0

4000

4200

4400

HeII 4686

40

Hγ

60

Hδ

Relative intensity

80

4600

4800

5000

5200

Wavelength (Å)

Fig. 4. Averaged, continuum-subtracted optical spectrum of J1305 obtained with SALT/RSS on the nights of 2012 April 15 - 16.

consistent with low-mass donors of spectral type K5 –M1 (Cox 2000).
As well as determining Porb from the Suzaku light
curves we can also attempt to place a limit on M1
by estimating K2 , the radial velocity semi-amplitudeof the donor, from the optical spectrum. Orosz et al.
(1994) and Orosz & Bailyn (1995) showed that there isa relation between the projected velocity of the outerdisc, vD , and K2 (vD /K2 = 1.1 − 1.25) where thepeak-to-peak separation of the Balmer lines can beused as an indicator of vD (Warner 1995). Utilizing adouble-Gaussian ﬁt to Hβ , we measure the peak-to-peakseparation to be 783 ± 37 km s−1 and derive a lower
limit to K2 ≥ 313 km s−1 .
Combining the preferred Porb with this conservative lower limit on K2 , we derive an absolute
lower limit on the mass function for J1305 off (M1 ) = M1 sin3 i/(1 + q)2 = K23 Porb /2πG ≥ 0.66M
where i is the binary inclination and q = M2 /M1 .
Since both sin i and q must be positive, f (M1 ) truly
provides an absolute lower limit on M1 . However,this is well below 1.4M, the canonical value for themass of a neutron star, and therefore does not allow
a determination of the true nature of the compact object.

4.3. Conclusions and Future Work
Using archival Suzaku data we have determined Porb of
the candidate BHXRT MAXI J1305-704 to be 4.98 ±0.05h based on the presence of deep dips in both the128s binned and folded Suzaku light curves. Though
this is the preferred period, we cannot completely ruleout other periods suggested by the Lomb-Scargle periodogram analysis. We also placed a lower limit on themass of the compact object M1 ≥ 0.66M , however thisdoes not yet constrain the nature of the compact object.
The work presented here is ongoing and requires further investigation of the optical spectrum. For example,
He ii λ4686 exhibits variability in the individual spectra,
perhaps indicating that at least some of the He ii emission originates on the X-ray heated face of the donor or
the stream/disc impact region, i.e. the hotspot. In order

to investigate this we plan to use Doppler tomography to
map the emission features in velocity space, recovering
the structure of the emission regions. Doppler tomography may also enable us to study the Bowen region in
greater detail and perhaps provide a better constraint
on K2 and therefore M1 . The best constraints on K2
will come from observations of the source in quiescence.
However, J1305 was not detected in the DSS, and therefore spectroscopy of the donor may require the next generation of optical telescopes such as the European Extremely Large Telescope (E-ELT).
References
Buckley, D. A. H et al. 2006, in IAU Symp., 232, 1
Charles, P. A., Coe, M. J. 2006, in Compact stellar Xray sources, Lewin W. H. G., van der Klis M., eds.,
Cambridge Univ. Press, p. 215
Charles P. A. et al. 2012, ATel, 4105,1
Cox, A. N. 2000, in Allen’s Astrophysical Quantities, 4thed., Springer, p. 389
Crawford, S. M. et al. 2010, in SPIE Conference SeriesVol. 7737, p. 25
Eggleton, P. P. 1983, ApJ, 268, 368
Greiner, J., Rau, A., Schady, P., 2012, ATel, 4030, 1
Kennea, J. A. et al. 2012a, ATel, 4034, 1
Kennea, J. A. et al. 2012b, ATel, 4044, 1
Matsuoka, M. et al. 2009, PASJ, 61, 999
Miller, J. M. et al. 2012, ATel, 4191, 1
Mitsuda, K. et al. 2007, PASJ, 59, 1
Orosz, J. A. et al. 1994, ApJ, 436, 848
Orosz, J. A., Bailyn, C. D. 1995, ApJ, 446, L59
Sato, R. et al. 2012, ATel, 4024, 1
Shidatsu, M. et al 2013, ApJ, 779, 26
Steeghs, D., Casares, J. 2002, ApJ, 568, 273
Tody, D. 1986, in SPIE Conference Series Vol. 627,
Instrumentation in astronomy VI. p. 733
Warner, B. 1995, in Camb. Astrophys. Ser., 28
Zurita, C. et al. 2008, MNRAS, 681, 1458

49

Bright novae– Indications of the spectrum of MAXI J0158−744 –Toshikazu Shigeyama,1 Kentaro Wada,1,2 Yukari Ohtani,31Research Center for the Early Universe, Graduate School of Science,
University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan2
Department of Astronomy, Graduate School of Science, University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan
3
National Astronomical Observatory, Mitaka-shi, Tokyo 181-8588, Japan
E-mail(TS): shigeyama@resceu.s.u-tokyo.ac.jp
Abstract
We are trying to construct a theoretical model describing bright novae with luminosities exceeding
1039 erg s−1 in the context of a steady wind model emanating from the surface of a white dwarf. The
motivation of this work originates from the discovery of a bright transient event MAXI J0158−744 in the
Small Magellanic Cloud. We succeeded in reproducing the peculiar spectrum showing a strong Ne IX
emission line at 0.92 keV with a large equivalent width and the absence of Ne X line at 1.02 keV andother emission lines from heavier ions like Mg and Al at higher energies by assuming a very fast wind
accelerated above the photosphere. This is in contrast with the existing nova wind model which assumes
acceleration below the photosphere. Thus we have extended the previous model to include the opticallythin region by utilizing a ﬂux-limited diﬀusion approximation. Though we can obtain solutions in which
the wind is accelerated above the photosphere, the luminosity cannot be so high as observed. We will
discuss possibilities to obtain more luminous solutions compatible with observed bright events including
MAXI J0158−744.
Key words: novae — stars: winds —line: proﬁles — scattering — radiation: dynamics

1.

MAXI J0158−744

A bright transient source MAXI J0158−744 was discovered by MAXI (Matsuoka et al. 2009) in the directionof the Small Magellanic Cloud on 2011 Nov. 11. Li etal. (2012) reported this discovery together with resultsof the follow-up observations by Swift and other opticaltelescopes. The detailed features of this object including the spectrum is given by Morii et al. (2013). The
continuum was ﬁtted by a black body spectrum with atemperature of kT ∼ 0.33 keV. This object became asbright as ∼ 1040 erg s−1 at maximum and lasted for∼ 10, 000 s. The inferred photospheric radius of a few
thousand kilometer suggests that the emission originatesfrom the surface of a white dwarf. Thus this transientphenomenon is a kind of nova explosions with an extremely short duration. Such a short duration indicatesthe mass of this white dwarf close to the Chandrasekhar
limit. Thus the estimated luminosity clearly exceeds the
Eddington limit.
The observed spectrum had a peculiar feature. Thereis a prominent emission line due to He-like Ne ions at 0.92
keV with an equivalent width of 0.32 +0.21−0.11 keV while no

 


Fig. 1. Schematic view of line blanketing eﬀects. See text for detail.

other lines were detected. This is strange because there
should be, for example, a strong emission line due to Hlike Ne ions with this observed photospheric temperatureof kT = 0.33 keV. We tried to account for this peculiar
spectrum by calculating radiative transfer of photons inan accelerating wind emanating from a white dwarf. In

50

1.1. Spectrum from accelerating nova wind
We have constructed a spherically symmetric steadywind model including optically thin regions. The velocity proﬁle v(r) is determined as a function of the radiusr by the momentum conservations:
dv
κL
v
=
dr
4πr2 c



LEdd
1−
L


, LEdd =

4πcGM
,
κ

(1)

where κ denotes the opacity, L the luminosity, c the
speed of light, LEdd the Eddington limit, G the gravitational constant, and M denotes the mass of the whitedwarf. The ﬁducial value for L is 8 × 1039 erg s−1 and M
is assumed to be 1.4 M . The solution of this equation
can be obtained as,

v=

κL
2πcR0


1−

LEdd
L


1−

R0
r


+ v02 ,

(2)

where R0 denotes the photospheric radius and v0 =v(R0 ). The density proﬁle ρ(r) is determined fromṀ = 4πr2 ρv,(3)where the constant mass loss rate Ṁ is determined toreproduce the observed photospheric radius of 2,300 km.
To determine the position of the photosphere, we calculate the ionization states of ions by the photo-ionizing
plasma package in XSTAR (Kallman & Bautista 2001).
The mass loss rate thus determined is 1.4 × 10−6 M
yr−1 , which is about 2 orders of magnitude smaller thanthat for a typical classical nova. The terminal velocity
becomes as high as 100,000 km s−1 . Note that this simple model does not specify the energy source and assumes
the luminosity of the observed value. This is diﬀerent
from the steady wind model presented in the next section, in which the obtained luminosity is required to becaused by the CNO-cycle in the vicinity of the surface ofa white dwarf.

   

 


expanding spherical matter with positive velocity gradients, a photon scattered oﬀ a certain ion reduces its
energy in the rest frames of any other ions at diﬀerent
places (Fig. 1). Thus a photon emitted by a certain
ion with a higher transition energy (Al XIII in Fig. 1)
is subsequently scattered oﬀ ions with lower transitionenergies (Mg and Ne ions). We expected that this line
blanketing eﬀects eliminate the other emission lines andhave succeeded in reproducing the spectrum. Furthermore we are able to extract some information of this phenomenon such as the duration and the elemental abundances (Ohtani et al. 2014). We summarize the resultsin the next section.

 

 















Fig. 2. Spectra with diﬀerent oxygen abundances. The oxygen abundance should be smaller than 5×10−9 to reproduce the prominent
emission line at 0.92 keV.

We calculate spectra by calculating radiative transferof photons emitted from the photosphere by the Monte
Carlo method (Ohtani et al. 2014). We ﬁnd that emission lines at energies higher than 0.92 keV are eliminatedby line blanketing eﬀects. The same eﬀect would eliminate the emission line at 0.92 keV if there were oxygenions with XO > 5 × 10−9 in the wind because OVIII has
a transition with an energy of 0.77 keV (see Fig. 2). The
CNO cycle converts most of O to N in 2,400 s after its termination. We argued that this could explain the strongline at 0.92 keV. The abundance of Ne to reproduce theobserved equivalent width needs to be XNe ∼ 0.01 or
more. Thus the white dwarf is composed of O and Ne
rather than C and O. In addition, the observed peculiar
spectrum suggests that the wind is accelerated above thephotosphere, which is at odds with the optically thickwind model Kato & Hachisu (1994) developed to reproduce light curves of novae till the beginning of the supersoft X-ray source phase. In this model, the sonic point
resides below the photosphere and the velocity becomesalmost constant before reaching the photosphere.
2. Steady wind model
To see whether there exist steady wind solutions in which
the wind is accelerated above the photosphere, we extend
the optically thick wind model to include the opticallythin regions by using the M1 closure method (Levermore
1984) in the calculation of the radiative energy ﬂux. Here
we present a summary of Wada & Shigeyama (2017) inwhich winds from CO white dwarfs are explored.
Governing equations are the conservations of the mass,

51

the momentum, and the energy given by,
4πr2 ρv = Ṁ ,
dp

GM ρ
κρF0
g
ρv dv
dr + dr = − r 2 +c ,1 d
r 2 dr

(4)
(5)





2
ρ
r2 v eg + ρv2 − GM
+
p
+
P
+
E
g
0
0
r
 2 
d
r F0 = ρ,+ r12 dr(6)

where pg the pressure due to ideal gas, G the gravitational constant, κ the opacity, F0 the radiative ﬂux inthe rest frame of matter, c the speed of light, eg denotes
the energy density of gas, P0 the radiation pressure, the energy generation rate due to nuclear reactions, E0
is the energy density of radiation. Here the subscript 0
indicates the quantity in the rest frame of matter.

Fig. 4. Evolutions of luminosity and temperature of the photospherein the HR diagram for white dwarfs with three diﬀerent masses.
Each ﬁlled square on a line represents the maximum luminosityand photospheric temperature attained by the corresponding optically thick wind model for each white dwarf.

Fig. 3.
Velocity (solid black line) and density (dashed redline) as functions of the radial coordinate for a nova windmodel with a white dwarf mass of M = 1.3 M , massloss rate of Ṁ = 3.895 × 10−7 M yr−1 , and luminosity of
L∞ = 2.275 × 1038 ergs−1 . The thin black lines denote escapevelocity (dotted) and sound speed (dot dashed) respectively. The
vertical dashed line indicates the radius of the critical sonic point
and the vertical solid line the photosphere.

The opacity κ is calculated using the opacity tablefrom OPAL (Iglesias & Rogers 1996) for a given chemicalcomposition. The chemical composition in the wind isassumed to be uniform and constant: We set the massfractions of hydrogen to XH = 0.35, helium XHe = 0.33,
heavy elements Z = 0.02, and the extra mass fractionsof CNO elements that are not included in Z is set todXCNO = 0.30.
The energy source of a nova is nuclear fusion reactions
on the surface of a white dwarf, primarily CNO-cycle
reactions. Instead of solving nuclear reaction network
equations, we use approximate formulae of the energy
generation rates for pp-chain, CNO-cycle, and triple alpha reactions as functions of the temperature T , themass density, and the mass fractions of elements (Kippenhahn et al. 2012). We have denoted the sum of thesethree rates as  in equation (6).
We numerically integrate the above three equations
from a point far beyond the photosphere (r = 109 km)
to the surface of a white dwarf with a given mass M
by the 4th order classical Runge-Kutta method with astep size of Δ ln r = 0.001. The boundary conditions are:i) a given luminosity at r = 109 km, ii) the luminosity
vanishing at the surface of the white dwarf. The othertwo are given by the mass loss rate Ṁ and the terminal
velocity v∞ (speciﬁed at r = 109 km). We use a shooting
method repeating the integrations with various pairs ofṀ and v∞ to seek a solution that smoothly passes thecritical sonic point and reaches the surface of the whitedwarf with a vanishing luminosity. Thus the mass loss
rate is obtained as an eigen-value to satisfy the boundary
conditions.

52

We have succeeded in obtaining wind solutions withvery small mass loss rates, which may describe the earlysuper-soft X-ray source phases of novae (Fig. 3). In
this phase, the bolometric luminosity of a nova becomes
higher than before, the sonic point moves very close to
the photosphere, and the wind is accelerated above thephotosphere. On the other hand, we are not able to obtain a solution with a luminosity as high as 1039 erg s−1even for a white dwarf with a mass close to the Chandrasekhar limit (see Fig. 4). Though there are transonicsolutions with such high luminosities, none of them satisﬁes the inner boundary conditions. The terminal velocities are a few thousand km s−1 at most (Fig. 5), which
is a factor of a few tens smaller than that required toreproduce the observed spectrum of MAXI J0158−744.

Fig. 5. Sequences of steady state solutions in the parameter spaceof the mass loss rate Ṁ and the terminal velocity v∞ . The thicksolid line (black), the thick dashed line (blue), and the thick dot
dashed line (magenta) denote models with white dwarf masses of
M = 1.0 M , 1.3 M , and 1.388 M , respectively.

3. Remaining problems
As shown in §1., the peculiar spectrum of MAXI
J0158−744 indicates that matter in the wind is accelerated above the photosphere. The positive gradient of
the velocity enables the line blanketing to eliminate emission lines at energies higher than 0.92 keV. On the other
hand, the paucity of O ions in the wind as a result of
the CNO cycle prevents the same mechanism from reducing the observed emission line at 0.92 keV. Since thecurrent optically thick wind model can not deal withwind solutions in which the matter is accelerated above
the photosphere, we have extended the model to include
the optically thin region by introducing the M1-closure
method. Though we obtain solutions in which the windis accelerated above the photosphere, the maximum luminosity is far (more than a factor of 10) below thatof MAXI J0158−744. As a result, the velocity and its
gradient in these steady models are not suﬃciently highfor line blanketing eﬀects to block photons with higher
energies.
The observed values of the photospheric velocities andluminosities together with the short duration of MAXI
J0158−744 indicate the mass of the white dwarf is close
to the Chandrasekhar limit. A massive white dwarf composed of O and Ne rather than C and O is likely to
reproduce the peculiar spectrum with a strong Ne emission line. With such a massive white dwarf, explosive
thermonuclear runaway of the CNO-cycle might lead toluminosities much higher than those attained by steadystate solutions investigated here. To explore such nonsteady behaviors, we may need to integrate hydrodynamical equations with respect to time for periods much
longer than the free-fall time scale at the surface of thewhite dwarf, which may need a signiﬁcant improvementin the existing hydrodynamical codes to simultaneously
trace explosive winds and nearly hydrostatic states of awhite dwarf for a long period. Otherwise, we need toproperly set boundary conditions at the surface of thewhite dwarf and calculate only the evolution of winds,
though we do not know appropriate boundary conditions.
References
Iglesias, C. A., & Rogers, F. J. 1996, ApJ, 464, 943
Kallman, T., & Bautista, M. 2001, ApJS, 133, 221
Kato, M., & Hachisu, I. 1994, ApJ, 437, 802
Kippenhahn, R., Weigert, A., & Weiss, A. 2012,
Stellar Structure and Evolution:
Astronomy
and Astrophysics Library. ISBN 978-3-642-302558. Springer-Verlag Berlin Heidelberg, 2012
Levermore, C. D. 1984, J. Quant. Spectrosc. Radiat.
Transfer, 31, 149
Li, K. L., Kong, A. K. H., Charles, P. A., et al. 2012,
ApJ, 761, 99
Morii, M., Tomida, H., Kimura, M., et al. 2013, ApJ,
779, 118
Matsuoka, M. et al. 2009, PASJ, 61, 999
Ohtani, Y., Morii, M., & Shigeyama, T. 2014, ApJ, 787,
165
Sugizaki, M. et al. 2011, PASJ, 63, 635
Wada, K., & Shigeyama, T. 2017, PASJ, submitted

53Classiﬁcation of Hard X-ray Light Curves of Black Hole Binary Transients
K. Ebisawa1 , S. Nahahira1 , T. Amagasa2 , F. Hayashi2 , Jun Fan2 , J. Sugimoto3 , A. Yoshida4
1Japan Aerospace Exploration Agency (JAXA), 2 Tsukuba University,
3
RIKEN, 4 Aoyama Gakuin University
E-mail(KE): ebisawa@isas.jaxa.jp
Abstract

We have systematically analyzed hard X-ray light curves of transient black hole binaries taken by
Swift/BAT (15−50 keV) and MAXI (3−20 keV), searching for common morphological properties among
diﬀerent light curves. To that end, we have applied a technique widely used in data science, such that
virtual “distance” between two lightcurves is deﬁned to quantify their similarities, and “tree diagram” ismade to classify all the observed lightcurves. Consequently, we have found that the lightcurves in 15−50
keV are classiﬁed into distinct groups based on their morphologies, while such a classiﬁcation is not clear
in the 3−20 keV light curves. This suggests that hard X-rays are more likely to reﬂect distinct types of
the accretion processes in black hole binaries than soft X-rays.
Key words: black holes: light curves – MAXI – BAT – Data Science

1. Introduction
Black hole transients exhibit varieties of X-ray light
curves in terms of their morphologies, presumably reﬂecting various aspects of diﬀerent accretion processes.
However, our current understanding of the physicalmechanisms behind these varieties is very limited.
Both Swift/BAT and MAXI routinely produce andpublicly release the X-ray light curves of bright blackhole transients. These light curve datasets will provide
a unique opportunity to carry out systematic and unbiased studies of the black hole X-ray transient phenomena. Here, we are going to apply techniques that are
proven eﬀective in various ﬁelds of “data sciences”, andtry to classify X-ray light curves of black hole transients
based on their morphologies.
2. Data Analysis and Results
We use the Swift/BAT (15–50 keV) and MAXI (3–20 keV) publicly available X-ray light curve datasets.
We choose the following black hole (candidate) sources;
4U1630–472, GX339–4, H1743–322, MAX J1305–704,
MAXI J1543–564, MAXI J1659–152, MAXI J1836–194,
MAXI J1910–057 and XTE J1752–223. For each pair of
all the BAT light curves, we calculate their Dynamic
Time Warping (DTW) distance, which is known to be aneﬀective measure to quantify their similarities (Hayashi
et al. 2013 ); we did the same for the MAXI lightcurves.
A merit of using the DTW distance is that it can ﬁndmorphological similarities between two light-curves even
if they have diﬀerent intensities or stretched in the time

direction. Based on these DTW distances, we performed
the clustering analysis to construct “tree diagrams” forBAT and MAXI light curves separately. On the tree diagram, light curves which are more similar in shape arelocated closer.
Consequently, the BAT tree diagram (Figure 1) shows
several distinct clusters, while MAXI tree diagram (notshown) does not show such clear clustering. This indicates that hard X-ray is more likely to reﬂect characteristics of accretion process than soft X-ray, whatever thephysical mechanisms behind. Looking at the BAT tree
diagram and those light curves belonging to each cluster,
we notice the following; besides a few “outliers”, ordinary light curves are divided into those having “doublepeaks” (Figure 2, top-left) or “single-peak”. The singlepeak light curves are further divided into “fast-decay”(top-right), “fast-rise” (bottom-right) and “slow-decay”(bottom-left) types.
Physical mechanisms behind these classiﬁcations areunclear yet. Also, it is curious why hard X-ray light
curves are more eﬀectively classiﬁed than soft ones. In
order to study their physical origins, we are investigating
for correlations between the light curve morphologicaltypes and black hole parameters, such as black hole mass,
binary parameters, luminosities and spectral states.
References
Hayashi, F. et al. 2013, Journal of Space Science
Informatics Japan: Vol.2, p.19

54

Fig. 1. A tree diagram to show classiﬁcation of the black hole light curves taken by Swift-BAT.

Fig. 2. Four representative examples of classiﬁcation of the Swift-BAT lightcurves.

55

Optical Monitoring of the MAXI sourcesby the optical telescope AROMA-N
Yoshiki Kitaoka,1 Takanori Sakamoto,1 Atsumasa Yoshida1

1Graduate School of Science and Engineering, Aoyama Gakuin University
5-10-1 Chuuou-ku Fuchinobe Sagamihara city Kanagawa pref 252-5258
E-mail(YK): ykitaoka@phys.aoyama.ac.jp
Abstract

MAXI is constantly monitoring variable X-ray sources and measuring the ﬂuxes in the X-ray bands.
Aoyama Gakuin University (AGU) has a small optical telescope called AROMA-N (AGU Robotic Optical
Monitor for Astrophysical objects-Narrow ﬁeld) which is dedicated to perform follow-up and coordinatedobservations of high energy astrophysical sources. AROMA-N is located on the roof of the building of
the AGU Sagamihara campus. This telescope is the Schmidt-Cassegrain type and 30 cm in diameter.
The german-type equatorial mount, the auto-focuser and the ﬁlter wheel with the Johnson UBVRI ﬁlters
are equipped. The typical limiting magnitude of AROMA-N is 16-17 mag in 60 second exposures. The
operation of AROMA-N is fully automated. Since AROMA-N constantly observes the MAXI sources, we
can track the ﬂux variations of the sources both in X-ray and optical bands combining the MAXI data.
In this paper, we show the results of MAXI sources observed both in X-ray and optical using the MAXI
and the AROMA-N data. We will report the analysis results and discuss the prospects for the future.
Key words: X-ray, Optical counterpart, Telescope1.

The details of AROMA-N and attached devices

We use the optical telescope AROMA-N (AGU Robotic
Optical Monitor for Astrophysical Objects-Narrow ﬁeld).
Fig 1-5 show the attached devices to AROMA-N.
Table 1: The conﬁguration of AROMA-N

Apeture305 mm

Focal length1470 mm

Telescope
Schmidt-Cassegrain

Fig. 1: AROMA-N

Equatorial
Garman type

Upper limit
16-17 magField of view
25’ × 25’

Fig. 2: CCD

Fig. 4: ﬁlter

Fig. 3: focuser

Fig. 5: reducer

2. The target objects of AROMA-N
AROMA-N is monitoring the sources in the MAXI/GSC
catalog (Hiroi et al. 2013). The catalog contains the
MAXI GSC sources detected in the 37 months data at
the high galactic-latitude sky. We are monitoring thepart of these objects by AROMA-N. AROMA-N detected the optical counterparts for 5 objects (table 2).
Table 2: MAXI/GSC catalog objects AROMA-N observed.
The color-hatched sources are detected by AROMA-N

Objects nameMAXI J0004+726
MAXI J0043+247
MAXI J0048+320
MAXI J0055+463
MAXI J0123+342
MAXI J0151+361
MAXI J1429+425
MAXI J1437+588
MAXI J1535+581
MAXI J1658+353
MAXI J1704+785
MAXI J1706+240
MAXI J1834+327
MAXI J2241+297
MAXI J2354+286

Type
X-ray source
Cluster of galaxies
Seyfert 2 Galaxy
Cluster of galaxies
BL Lac object
Cluster of galaxies
BL Lac object
Seyfert 1 Galaxy
Seyfert 1 Galaxy
Low mass X-ray binary
Cluster of galaxies
Semi-regular pulsacng star
Seyfert 1 Galaxy
X-ray source
X-ray source

RA,DEC(00:04:48.00,+72:42:00.00)(00:43:52.00,+24:24:21.00)(00:48:47.14,+31:57:25.08)(00:55:19.90,+46:12:57.00)(01:23:08.65,+34:20:48.56)(01:52:50.40,+36:08:46.00)(14:28:32.61,+42:40:21.05)(14:36:22.13,+58:47:38.93)(16:28:37.00,+39:31:28.00)(16:57:49.81,+35:20:32.40)(17:03:43.50,+78:43:03.00)(17:06:34.52,+23:58:18.55)(18:35:03.39,+32:41:46.86)(22:42:00.00,+29:42:00.00)(23:54:24.00,+28:36:00.00)

56

3.

The results of MAXI/GSC catalog objects observed by
AROMA-N
AROMA-N detected the optical counterparts for 5 objects listed in table 2. We analyzed I, R, and B band
images, and extracted the source magnitudes. (table 3)
Table 3: The magnitudes of detected MAXI/GSC sources by AROMAN

Band
Magnitude
Band
Magnitude
Band
Magnitude

2016 7/6 MAXI J1437+588
I bandR bandB band12.46 (0.025) 13.00 (0.036) 14.82 (0.063)
2016 11/3 MAXI J1535+581
I bandR bandB band13.32 (0.032) 14.04 (0.062) 15.57 (0.127)
2016 7/6 MAXI J1658+353
I bandR bandB band13.50 (0.041) 13.30 (0.039) 14.29 (0.037)

4. Comparison with X-ray and optical bands
We monitored the optical counterpart of MAXI
J0055+463 and MAXI J1834+327 from 2016 10/14 to2016 11/9 by AROMA-N. We compared the light curves
in X-ray (2-20 keV) and optical (I, R, B bands). Figure
6 show the light curves of MAXI J0055+463 and MAXI
J1834+327.

5. New X-ray source MAXI J0636+146
In addition to the MAXI/GSC catalog sources,
AROMA-N is monitoring new X-ray sources detectedby MAXI. AROMA-N observed the new X-ray source,
MAXIJ0636+146 , which was detected by MAXI/GSC
at 02:20 (UT) on 2016 November 4 (Negoro et al. 2016).
Three days after the discovery, AROMA-N collected several images of the ﬁeld. However, we did not detect the
optical counterpart at the Swift/XRT position (2.5 arcsec error region).
Table 4: AROMA-N observation of MAXI J0636+146

RA,DECBand
5σ upper limit(06:37:24.18,+14:37:53.20)
R bandB band17.0
17.1

Fig. 7: R band image of MAXI J0636+146

Fig. 8: B band image of MAXI J0636+146

6. Summary and future
We monitored about ∼20 MAXI/GSC catalog sources
and compared X-ray and optical data. We detected theoptical sources for 5 objects. In the future, we will
expand the number of monitoring sources and further
investigate the relationship between X-ray and opticalbands.
References
Hiroi et al. 2013, ApJS, 207, 36
Negoro et al. 2016, ATel.9707, 1N

Fig. 6: X-ray and optical light curves of MAXI J0055+463, MAXI
J1834+327

57

A panchromatic overview of accreting binary systems and their associatedrelativistic jets
Stephane Corbel1
1

Univ. Paris Diderot & CEA Saclay & Observatoire de Paris
E-mail(SC): stephane.corbel@cea.fr
Abstract

In the past decade, several considerable achievements have been reached in the ﬁeld of Galactic accreting
binary systems, especially in light of the extreme variability of their relativistic jets. These jets have moved
from being considered exotic and rare abnormalities to being recognized as integral and vital components
in the transfer of energy and angular momentum.
Although their phenomenology is now rather well established, their emission and contribution to thetotal energy budget of microquasars at large (and connexion to the supermassive black holes) is still thesubject of active debates. I will present the most relevant observations concerning our understanding ofrelativistic jets in accreting systems, discussing in particular some open issues.

58

59

Search for soft X-ray ﬂashes at the ignition phase of optical novae using
GSC data
Mikio Morii1
1

Institute of Statistical Mathematics
E-mail(MM): morii@ism.ac.jp
Abstract

We searched for soft X-ray ﬂashes associated with optically discovered classical or recurrent novae in
5-yr MAXI/GSC data in 2–4 keV band. For this purpose, we developed a tool to make the light curvesof a point source by ﬁtting the event distribution with point spread functions (PSF-ﬁt tool). We applied
the PSF-ﬁt tool to 40 classical/recurrent novae that were discovered in optical observations from 2009
August to 2014 August. We found no precursive soft X-ray ﬂashes with signiﬁcance above the 3-sigma
level between T − 10 day and T (T = discovery date). We obtained the upper limits for the bolometricluminosity of soft X-ray ﬂashes, and compared them with the theoretical prediction and that observed for
MAXI J0158−744.
Reference: M. Morii et al. 2016, PASJ, 68, S11

60

61

Stellar Mass
Black Holes

62

63

Recent Constraints on Jet physics and Properties Obtained from HighEnergy Observations of Microquasars
Jérôme Rodriguez1
1

Lab AIM - CEA/IRFU-CNRS/INSU-Université Paris Diderot,
CEA DRF/IRFU/SAp, F-91191 Gif-sur-Yvette, France.
E-mail: jrodriguez@cea.fr
Abstract

The recent detections of microquasars at energies above a few hundred keV up to the TeV in one casehas stimulated a strong interest and raised several questions. How are the MeV, GeV, and even TeVemissions, produced? What are the emission processes and in-ﬁne the physics and media at the origin of
the broad band spectra? What is the content of these media, and how are they powered? These sources
are machines accelerating particles and matter to very high speed in collimated jets. While these are now
known for more than twenty years (mainly with radio observations) and their behavior known to be tied tothe accretion processes onto the compact object (studied via X-ray observations), their potential inﬂuenceat high energy is just being recognized. In this review I will present a selection of recent results obtainedwith XMM-Newton, INTEGRAL, Fermi on 4U 1630−40, Cygnus X-1, and V404 Cygni and will discuss
the possible interpretations of these results.
Key words: accretion, accretion disks — jets sources — black hole physics — X-rays: binaries —Gamma-rays: observations — stars: individual: Cyg X-1, V404 Cyg, 4U 1630−47

1.

Introduction

Microquasars (MQs) are Galactic X-ray binaries showing episodes of ejections in certain phases of their activity. Most MQs host a black hole as compact object, andthe majority are transient sources that spend most of
their lives in a dormant, so-called quiescent state. They
are detected (usually in X-rays) when they enter intomonth to year long periods of activity called outbursts.
All MQs are multi-wavelength emitters. The radio to
infrared (IR) emission is attributed to relativistic jets,
either a persistent compact jet or in form of large scale
discrete ejections. The X-rays probe the inner accretion
ﬂow(s): the accretion disc (∼0.1–10 keV) and the socalled “corona” (10–200 keV) (e.g. Fender, 2006; Remillard & McClintock, 2006). The relative contributionsof these two media seen in the X-ray spectra have led
to the classical classiﬁcation into spectral states. The
two canonical ones are the “(high) soft state” (HSS) andthe “(low) hard state” (LHS). In the former the X-ray
spectrum is dominated by the thermal emission from thedisc, and in the latter it is dominated by a (cut-oﬀ) powerlaw with a hard (Γ<2) photon index attributed to inverse
Comptonization (IC) of the soft disc photons by the coronal electrons. The past 25 years have allowed us to enrich
this initial vision, and, in particular, to ﬁnd strong con-nections between X-rays (accretion) and radio-IR (ejection) (e.g. Mirabel et al., 1998; Fender & Pooley, 1998;
Corbel et al., 2000; Fender, 2001b; Corbel et al., 2003;
Rodriguez et al., 2003, 2008; Corbel et al., 2013b, to cite
but a few): a powerful compact jet is always present in
the LHS, while transitions from the LHS to the HSS areaccompanied by discrete ejections.
The picture may have recently been completed by
the detection of a few sources at energies above 1 MeV
with OSSE and INTEGRAL, and/or in the GeV domain
with Fermi and Agile, and possibly in the TeV range
with the Magic telescope. These detections pose questions regarding the media and physical processes at work,with answers that will certainly impact our understanding of particles acceleration, particle-matter and particleparticle interaction, and eventually will help us constrain
feedback processes in the interstellar medium.
As mentioned above these sources are either transients, and they are anyway highly variable. One thus
needs alerts from surveying instruments in order to beable to conduct multi-wavelength, multi-instrumental
campaigns during the most relevant periods of their activity. Before presenting recent high energy observations
and analysis of a few MQs, I will ﬁrst start this review
by presenting the obvious but necessary needs of all sky

64

observing campaigns during speciﬁc periods of activity.

"Failed"outburst

Spectral transitions

Fig. 1. MAXI light curves of GX 339−4 obtained from the MAXI
online data at http://maxi.riken.jp/top/

surveying/monitoring instruments without which mostof the results of the past 30 years will not have beenobtained.
2.

Transient and variable sources need to be surveyed: thecentral role of all-sky monitors
MQs have unpredictable behavior. They enter into outburst and/or show (spectral) transitions at “random”
moments. These are ideal times to probe the interactions and causal relations between the diﬀerent media(jet-disc-corona-companion). It is thus necessary to keep
an eye on the entire sky to detect new or know object asearly as possible during their outbursts, or simply checkfor any spectral transition that may ’act’ as a smoking
gun for major change of conﬁguration of the inner ﬂows
and possible ejection events. Wide ﬁeld monitoring instruments (hereafter ASM for all sky monitors of any
types) are the only instruments to do so.
Fig. 1 shows the MAXI light curves covering a multiyear survey of GX 339−4 that, in many ways, can be considered as the prototype MQ (e.g. Corbel et al., 2013a).
Over the 7 years of X-ray coverage, one can clearly distinguish three outbursts separated by long intervals ofquiescence. The outbursts, moreover, show very diﬀerent
proﬁles, with two long (>6 months), bright (>100 mCrab
at all energies) episodes, and a shorter and much fainterone (dubbed “failed” in Fig 1). The spectral capabilities of ASMs permit to have an overview of the spectral
dependence of the sources light curves, build hardness ratios and hardness intensity diagrams (HID). This allows
to survey the sources spectral evolution/transitions andclassify their states (Fig 2) to eventually trigger/organise

It is thus possible to use the entire database of any
ASM to classify a source’s behavior from multi-color
light-curves and HID without a full model dependent
approach to the data (e.g Punsly & Rodriguez, 2013;
Grinberg et al., 2014, in the respective cases of GRS
1915+105 and Cyg X-1; see Fig. 2). This also allows
to ﬁll the temporal gaps and/or uneven temporal sampling of dedicated pointed observing campaign madewith more eﬃcient, smaller ﬁeld instruments (XMMNewton, Chandra, RXTE, Nustar, for instance). Thisis an approach that we followed and that allowed us toreﬁne the study of the high energy properties of Cyg X-1(Rodriguez et al., 2015b).
3. Cyg X-1 and the possible origin hard spectral tails seenin (some) microquasars
While emission above 100 keV from microquasars is
known for more than 20 years (e.g. Laurent et al., 1993;
Churazov et al., 1994; Grove et al., 1998; McConnell
et al., 2000, to cite just a few), no consensus on its origin has been reached yet. It can be interpreted as due
to IC, either purely thermal (spectrum with an exponential cut-oﬀ at about 100 keV, Titarchuk, 1994; Rodriguez
et al., 2003; Cadolle Bel et al., 2007), or thermal/nonthermal in an hybrid plasma (e.g Zdziarski et al., 2001;
Del Santo et al., 2013), or from thermal IC from two different populations of electrons (e.g. Bouchet et al., 2009,
to ﬁt the 1E 1740.7−2942 INTEGRAL spectrum). In
the LHS, jet emission can also be invoked (e.g Markoﬀ
et al., 2005) through various radiation processes (Comptonisation, Synchrotron and Synchrotron-Self Compton- SSC). Apart from pure thermal IC, all these models
predict similar 20–1000 keV spectral shape, and a single spectral analysis, even with the high quality dataobtained today, does not permit to break the model degeneracy (see also Nowak et al., 2011).
Using the large Cyg X-1 INTEGRAL archive Laurent
et al. (2011) with IBIS, and Jourdain et al. (2012) with
SPI , reported the detection of a high energy tail extending to about 1 MeV in addition to a standard IC
component below typically 200–300 keV, and a strong
polarized signal above ∼400 keV. The large polarisation
fraction reported ( 60%) led both teams to conclude
that the origin of this > 400 keV tail was Synchrotron
emission coming from a compact jet (Laurent et al., 2011;
Jourdain et al., 2012).
In both studies, however, the entire INTEGRAL dataacquired before 2009 was used, mixing thus all spectral
states. As no radio data was considered by any of thetwo teams the presence of a jet could not be proved either although Cyg X-1 is a known jet source. To go
further, and discriminate the potential state dependence

65

Fig. 2. Left: Flux-Hardness diagram based on the RXTE/ASM data. The diﬀerent regions show diﬀerent spectral states as conﬁrmed witha spectral analysis. Right: RXTE/ASM, MAXI, and Ryle/AMI 15 GHz light curves of Cygnus X-1 over 16 years. Adapted from Grinberg
et al. (2014).

Ph/cm2/s

Pf=75±32%
=40°±13°

Pf<25%

Fig. 3. LHS INTEGRAL spectrum of Cyg X-1. The shaded parts indicate the spectral range were polarization measures were obtained,with the levels obtained in each cases. The red (resp. green) line
represents the IC (resp. power law) component. Adapted from
Rodriguez et al. (2015b).

of the results, we classiﬁed the whole INTEGRAL and
AMI-Ryle databases into three basic states based on the
ASM-based model independent classiﬁcation (Grinberg
et al., 2014, in Fig. 2, Sec. 2.). Stacked high energy spectra, polarigrams and radio properties were then obtained
and studied for all three states. A high energy tail, dominating the spectrum above ∼400 keV, is visible only in
the LHS (Fig. 3 Rodriguez et al., 2015b). A high degree(∼ 75%) of polarisation is also measured in this state atenergies where the tail dominates, while only an upper
limit is obtained at lower energies, where the spectrum
is dominated by IC, and in the other states. The radio observations independently indicate a high level ofemission (a mean of about 13 mJy) in the LHS while in

the HSS we report a mean level of 3 mJy, probably due
to residual emission from discrete ejections occurring at
LHS to HSS transitions. Our reﬁned analysis therefore
clearly conﬁrms the presence of both a highly polarized>400 keV tail and a compact radio jet during the LHS.
We interpret this tail as Synchrotron emission from thejet (Laurent et al., 2011; Jourdain et al., 2012; Rodriguez
et al., 2015b).
Alternative possibilities, such as hybrid thermal-nonthermal Comptonization in the corona (Zdziarski et al.,2014), or Synchrotron from the corona (see Romero
et al., 2014, for all details), have, however, also been suggested. The clear detection of Cyg X-1 with Fermi/LAT
up to ∼10 GeV (Bodaghee et al., 2013; Zanin et al., 2016;
Zdziarski et al., 2016) in the LHS, in addition to the possible TeV detection with MAGIC during a ﬂare (Albert
et al., 2007) renders the non-jet interpretation less likely(Zdziarski et al., 2014; Pepe et al., 2015). In particular,
and as I discuss in Sec. 5., the coronal model implies
a conﬁguration that is not entirely clear, and the MeVGeV emission is better modeled by emission from a jet(Zdziarski et al., 2014; Pepe et al., 2015; Zanin et al.,2016; Zdziarski et al., 2016; Zhang et al., 2017, Sec. 5.),than by a model invoking IC up to the MeV.

4. The extreme case of V404 Cygni
The low mass microquasar V404 Cygni entered into outburst on June 15, 2015 after more than 25 years of quiescence. It was detected with all X-ray wide ﬁeld instruments (Barthelmy et al., 2015; Negoro et al., 2015;
Kuulkers et al., 2015). It was soon recognized that the

66(Rodriguez et al., 2015a; Siegert et al., 2016; Loh et al.,2016; Muñoz-Darias et al., 2016). Loh et al. (2016) report a ∼4σ detection in the 0.1–100 GeV Fermi observations at a position coincident with V 404 Cyg. The
detected emission at GeV energies occurs in conjunction
with the brightest hard X-ray ﬂare, itself associated witha spectral change in the source. Both events appear soonafter (6 hr) a giant radio ﬂare was detected (Loh et al.,2016). The temporal and spatial coincidence of all events
strongly support a γ-ray ﬂare from V404 Cyg during amajor ﬂare. The simultaneous detection of the 511 keVannihilation line and γ-rays of higher energies imply thatthe latter emission occurs far from the corona, pointing
in the direction of jet emission (Loh et al., 2016). The
multi-wavelengths properties near the Fermi detection
are actually compatible with a Blandford-Znajek (BZ)
powered jets interpretation during states when the inneraccretion disk is magnetically arrested (Loh et al., 2016).







  

source had a dramatic behavior, showing intense ﬂares(up to 50 Crab in the 20–40 keV range) on timescalesof about an hour (e.g. Rodriguez et al., 2015a; Roques
et al., 2015a; Jourdain et al., 2017, Fig. 4). It was also
clearly detected up to high energy early in the outburst(Roques et al., 2015b), and a multi-instrumental followup and analysis showed probable detection of the sourcewith Fermi in the GeV range (Loh et al., 2016, Fig. 4).
Until then all GeV binaries were high mass systems
with either a Be or a Supergiant companion. In thesesystems, the GeV emission is explained by IC scatteringof the stellar (UV) photons on relativistic electrons near
the compact object (this model is known as the leptonic
model, e.g Bednarek, 2013; Dubus, 2013, and references
therein, Sec. 5.). In V404 Cyg, however, the secondary is
a 0.7+0.3−0.2 M K3 III companion (Khargharia et al., 2010,e.g.). In the context of the leptonic model, one cannot
expect GeV emission from this source. The discovery of
a variable transient feature in the 511 keV region, attributed to e− /e+ annihilation further brings the question about the production of pair plasma (Siegert et al.,2016), and the potential role of a jet (that is detected atother wavelengths). The annihilation ﬂux seems correlated with the X-ray ﬂaring activity, which itself somehow correlates with the ﬂares seen at longer wavelengths







Fig. 4. Multiwavelength light-curve of V404 Cyg during the June2015 outburst. From top to bottom AAVSO, Swift BAT, INTEGRAL/ISGRI and INTEGRAL SPI 511 keV, Fermi/LAT light
curves, and statistical test of the Fermi data. Adapted from Loh
et al. (2016); Siegert et al. (2016), courtesy A. Loh.

5. Models for high energy emission and jet contents
The origin of the high energy (>MeV) emission is debated. Two main families of models co-exist and predict
strong diﬀerences at very high energies only (TeV), a
domain were there is, to date, only one marginally signiﬁcant detection of Cyg X-1 (Albert et al., 2007). In
the ﬁrst family (leptonic models) the high energy emission is due to IC scattering of the strong stellar photon ﬁeld (until 2015 all GeV detected binaries were highmass systems) on relativistic electrons. The second family (hadronic models) invoke the presence of a jet, andthe MeV-GeV is instead due to post break Synchrotron
and SSC radiations (see e.g. Pepe et al., 2015, and references therein). Leptonic models are popular in thecontext of γ-ray binaries where the companion is a highmass star and the primary a pulsar.
In the LHS of MQs the radio-IR emission (Fig. 5) is well
!"#$!

!"%!"  "! !"Disc "Corona"

Jet











Jet



















Fig. 5. Lepto-hadronic modeling of the Cyg X-1 broad band spectrumwith a set of parameters favoring Synchrotron emission and SSC
from the jet as the origin of the MeV-GeV component. Adaptedfrom Pepe et al. (2015).

67

explained by self-absorbed Synchrotron radiation froma population of electrons in a compact jet (e.g. Blandford & Konigl, 1979; Fender, 2001b, 2006; Corbel et al.,2013b). Here the presence of the relativistically accelerated particles in a strongly collimated outﬂow mightchange our interpretation of the origin the high energy
emission(s). The strong polarization fraction detected inthe γ-ray tail of Cyg X-1 excludes IC (Rodriguez et al.,2015b), but not necessarily coronal emission (Romero
et al., 2014). However the latter explanation contains
intrinsic diﬃculties that led us to favor the jet emission. The high level of polarization measured with INTEGRAL needs an highly ordered magnetic ﬁeld which islikely to be found close to the base of the jet, where highenergy emission could also occur1 . Instead, Synchrotron
photons produced by coronal electrons would either haveto undergo many Compton scattering if produced deep
into the corona, and the polarisation would be lost, orthey would be produced too far from the highly orderedﬁeld to explain the high degree of polarization we measured.
Pepe et al. (2015) recently applied a lepto-hadronicmodel to represent the broad band spectrum of Cyg X-1,
the hadronic component being required by the TeV ﬂaredetected by MAGIC. While globally, the two sets of parameters they use ﬁt the data rather well, the one basedon coronal IC and IC of the companion star fails to represent the MeV tail, while this component is well ﬁtted
when considering direct jet Synchrotron and SSC instead(Fig. 5). The recent detection of MeV and GeV emission during the V404 Cyg outburst (Siegert et al., 2016;
Loh et al., 2016; Jourdain et al., 2017) implies similar
conclusions, especially in a system where the secondary
has no wind and a much weaker and cooler photon ﬁeld
than high mass systems.
Knowing the composition of jets is of prime importance as it permits to have a proper estimate of their
total kinetic power and in-ﬁne their feedback (matterand energy) into the ISM. Baryonic jets should indeed
cary signiﬁcantly more energy than pure leptonic ones,
and this of course puts even stronger constraints on themechanisms responsible for their production and acceleration. It seems widely admitted that jet might carya very signiﬁcant part of the accretion energy (Fender,
2001a; Gallo et al., 2005), in the LHS. Jet acceleration isthought to either be based on the extraction of the BH
spin (BZ) or magnetic energy from the disc (Blandord &Payne mechanism). Until recently, SS 433 was the only
MQ were baryonic jets had been detected (e.g. Kotani
et al., 1994), but this source is a peculiar MQ in manyother ways, and this could be just another of its peculiar*1 It is interesting to note that Zhang et al. (2017) could also
reproduce the polarized hard tail with a jet model involving
both large and small scale turbulent magnetic ﬁelds.

ities. In this respect the recent XMM-Newton detectionof Doppler shifted emission lines in 4U 1630−47 (Dı́az
Trigo et al., 2013), a more typical MQ, may shed newlight on this aspect. Dı́az Trigo et al. (2013) could use asimultaneous coverage at radio wavelengths that showedthat the Doppler shifted lines are coincident with the detection of a jet. The spectral analysis of the X-ray lines
are compatible with being emitted from a ∼0.7 c outﬂow. The ﬂux ratio between the red and the blue shifted
lines is, furthermore, consistent with Doppler boosting(Dı́az Trigo et al., 2013). This exciting result may imply
that typical MQ jets do carry baryons, and that, therefore, their energetic budget is even more important thanpreviously thought. Additionally, and as discussed by
Dı́az Trigo et al. (2013), the detection of a baryonic jet
favours models of accretion disk powered jet rather than
BZ jets. A baryonic jet implies that 4U 1630−47 should
be a strong source of γ-rays, and this source is not a
known as an high energy emitter. These lines have, furthermore never been conﬁrmed since their ﬁrst reportwith new Chandra or re-analysis of the XMM-Newton
observation (Neilsen et al., 2014; Wang & Méndez, 2016).
It is therefore not possible to fully conclude on this issue,
which leaves the mystery complete.
6. Conclusions
Jets are clearly important channels for energetic, material redistribution and feedback into the ISM. Theirdetection at high energy will certainly modify our understanding of MQs in general. The chosen results described in this short review are just the tip of the iceberg,
and certainly do not allow to have ﬁrm conclusions onmost of the questions I raised in the introduction. Soon,the advent the CTA at TeV energy, and/or the detection of neutrinos from a MQ, in conjunction with observations at all wavelengths may help to make a step forward into many of these aspects. One should still keep inmind that studying the accretion-ejection physics relies
on multi-instrumental campaigns triggered at appropriate times. This cannot be done without constant surveyof the transient sky, and our ability to understand better
these objects deﬁnitely rely on all sky monitorings suchas those made with RXTE/ASM in the past, or currently
with MAXI, and Swift.
Acknowledgements
I warmly thank C. Gouiﬀès, P. Laurent, & A. Loh for useful discussions and comments on this manuscript. Partial
funding support is provided through the CHAOS project(ANR-12-BS05-0009), and UnivEarthS Labex (ANR-10LABX-0023 and ANR-11-IDEX-0005-02), and by the
French Space Agency (CNES) in support to the analysis of the INTEGRAL data.

68

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69

Violent Optical Variations and Correlation between Optical and X-ray
Variability During the 2015 Outbursts in V404 Cygni
Mariko Kimura,1 Keisuke Isogai,1 Taichi Kato,1 Yoshihiro Ueda,1 Satoshi Nakahira,2 Megumi Shidatsu,3Teruaki Enoto,1,4 Takafumi Hori,1 , Daisaku Nogami,1 Akira Imada1 , Colin Littleﬁeld,5 Ryoko Ishioka,6Elena P. Pavlenko,7 Alexei S. Pozanenko,8,9 , Hyungsuk Tak10 and the Variable Star Network (VSNET) team1

Department of Astronomy, Graduate School of Science, Kyoto University, Oiwakecho, Kitashirakawa,
Sakyo-ku, Kyoto 606-8502, Japan2
JEM Mission Operations and Integration Center, Human Spaceﬂight Technology Directorate, Japan
Aerospace Exploration Agency, 2-1-1 Sengen, Tsukuba, Ibaraki 305-8505, Japan
3
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
4
The Hakubi Center for Advanced Research, Kyoto University, Kyoto 606-8302, Japan
5
Astronomy Department, Wesleyan University, Middletown, CT 06459 USA
6
Institute of Astronomy and Astrophysics, Academia Sinica, 11F of Astronomy-Mathematics Building,
AS/NTU. No.1, Sec. 4, Roosevelt Rd, Taipei 10617, Taiwan (R.O.C.)
7
Crimean Astrophysical Observatory, 298409, Nauchny, Republic of Crimea
8
Space Research Institute, Russian Academy of Sciences, 117997 Moscow, Russia
9
National Research Nuclear University MEPhI (Moscow Engineering Physics Institute), Moscow, Russia
10
Statistical and Applied Mathematical Sciences Institute, Durham, NC, USA
E-mail(FS): mkimura@kusastro.kyoto-u.ac.jp
Abstract
We performed the optical photometry of the two outbursts in 2015–2016 in V404 Cyg, a black-hole
X-ray transient and investigated the connection between our optical data and the X-ray data derived from
Swift/XRT, Swift/BAT and INTEGRAL IBIS/ISGRI monitoring. Analyzing the simultaneous opticaland X-ray data, we found large-amplitude optical variations with regular patterns (the period: 30 min–2.5
hours, the amplitude: about 2 mag) at low luminosity (even at one hundredth of the Eddington luminosity)in the June outburst for the ﬁrst time in black-hole X-ray transients. Prior to the outburst, repetitive
variations in these objects had been observed only at high luminosity (e.g., GRS 1915+105) and the existing
theories were developing to explain these observations before the outburst. The correlation between therepetitive optical variations and the X-ray ones was good. With the time lag between them and the multiwavelengths SED analysis, we demonstrated that the repetitive optical variability was mainly dominatedby reprocessing of X-ray irradiation in the disk (Kimura et al. 2016). Also in the December outburst, therepetitive optical variations were detected in one term. In addition, we found that the X-ray emission
delayed to the optical one by about 30–50 s. The time lag inconsistent with both disk reprocessing andsynchrotron emission related to jet ejections. There is a possibility that many kinds of optical variations
having diﬀerent origins were observed during the two outbursts in 2015 in V404 Cyg.
Key words: accretion, accretion disc – black holes physics – binaries: general – X-ray: stars – stars:
individual (V404 Cygni)

1.

Introduction

V404 Cyg is one of transient black-hole (BH) low-mass
X-ray binaries (LMXBs). LMXBs are composed of a
black hole or a neutron star (the primary) and a lowmass star (the secondary) (Lewin et al. 1995). They are
known to exhibit sporadic outbursts, which are causedby the thermal instability due to partial ionization ofhydrogen (see e.g., chapter 5 Kato et al. (2008)). Thissystem underwent outbursts in June and December in2015 after the 26 years dormancy.
2. Observations
We performed the CCD photometry in the optical IC ,RC , V and B bands and with the clear ﬁlters with tele-

70

scopes having 20 cm–2 m diameters at 27 and 17 sites
during the June and December outbursts, respectively.
The observers measured magnitudes of V404 Cyg relative to local comparison stars whose magnitudes were
measured by A. Henden from the AAVSO Variable Star
Database1 or taken from United States Naval Observatory (USNO) B1.0 catalogue or USNO A1.0 catalogue
or the fourth U.S. Naval Observatory CCD Astrograph
Catalog (UCAC4).
Iíband − 3 mag

Bíband + 6 mag

RíbandVíband + 3 mag

Swift/BAT 15í50 keV × 106 counts
Swift/BAT 50í150 keV × 104 counts

INTEGRAL/IBIS ISGRI 25í60 keV
Ic bandRc band + 2 mag
V band + 5 mag
No filter + 8 mag

1013
1012
10

10−1

5
10−2

1011

(b)

10

9

10

8

107

20

106

10−3
10
10−4
count sec

Magnitude

1010

10−5
15
10−6

105

10−7

104
103

L bol L Edd

a

3. Results
3.1. Overall Light Curves
The overall optical and X-ray light curves in the 2015
June and December outbursts are displayed in Figures
1 and 2, respectively. We found the whole trend of the
optical light curves, a slow rise and rapid decay, resemble
the X-ray ones in both two outbursts. The duration ofthese outbursts were a few weeks. The amplitude of the
December outburst was smaller than that of the June
outburst. The interval between the two outbursts wasless than 6 months.

20
10−8

30
102

10−9

101
100
0

Magnitude

b

2

4

6
8
10
12
BJDí2457189 (day)

14

16

18

25
10−10
0

2

4

6

8

10

12

14

12
RíbandNo filter

1 hour13
14
15

first detection by Swift/BAT
16
0.30

0.35
0.40
BJDí2457189 (day)

0.45

0.50

Fig. 1. Optical and X-ray light curves of V404 Cyg during an outburst in 2015 June–July. Panel a shows overall multi-colour light
curves and Swift/BAT light curves. The plotted points are averaged for every 0.67 days. Panel b is an enlarged view of theshaded box in panel a (the ﬁrst detection of short-term variations). On BJD 2,457,203, the mean magnitude dropped below
V =17.0. Superimposed on this rapid fading, the amplitude ofvariations became progressively smaller and smaller. After BJD
2,457,205, the mean magnitude seemed to be constant, and theoutburst virtually ended.

*1 <http://www.aavso.org/vsp>

Fig. 2. Overall light curves in the IC , RC , V bands and with no ﬁlter
and in the X-ray 25–60 keV energy band during the outburst ofV404 Cyg from December, 2015 to January, 2016. The horizontal
axis represents days from BJD 2457380. The plotted open circles
are averaged magnitude for every 1 day. Here, LEdd is equal to1.35 × 1039 [erg/s].

3.2. Correlation between Optical and X-ray Variations
It is noteworthy that V404 Cyg exhibited violent optical variations with regular patterns throughout the June
outburst. Some examples are shown in Figure 3. We
classiﬁed them into two types. One is ‘dip-type’ oscillations which consist of repetitions of a gradual rise
followed by a sudden dip of ∼2 mag amplitude, sometimes with accompanying spikes on timescales of ∼30
min to ∼2.5 hours. The other is ‘heartbeat-type’ oscillations which have rhythmic small spikes of ∼0.1 mag
amplitude with short periods of ∼5 min. The dip-type
oscillations arose not only at high luminosity near the

71
104

4
10

10

(a)

11

30 min

103

6

11
12

102

8

12
13
13
14
10.44

10.46

10.48

10.50

(b)

10

30 min
11.36

11.38

11.40

11.42

101

12
100

14

6.4

6.5

6.6

6.7

6.8

Fig. 3. Repetitive optical variations during (a) BJD 2457199.435–2457199.52 and (b) BJD 2457200.34–2457200.44 in the 2015 June/Julyoutburst and the correlation between repetitive optical and X-ray variations (the right panel). (a) an example of dip-type oscillations. (b):an example of heartbeat-type oscillations. The blue, green, red, pink and navy dots represent the light curves in the B, V , RC , IC bands
and in the INTEGRAL 25–60 keV X-ray energy band.

Eddington luminosity (LEdd ) but especially at low luminosity (∼0.01LEdd ) in the outburst. The bolometricluminosity was estimated from the Swift/BAT survey
data in the 15–50 keV energy band and INTEGRALIBIS/ISGRI monitoring data in the 25–60 keV energy
band. This kind of oscillations were also seen for one
interval in the December outburst.
Moreover, the same kind of variations were observedin optical and X-ray wavelengths when the dip-type oscillations were observed (see the right panel of Figure3). The time delay of optical variations to X-ray ones
were estimated to be about 1 min by CCFs. In addition,
the estimated disk radius was ∼2.0 ×1012 [cm] with themulti-wavelengths SED analyses. Additionally, it wasdemonstrated the synchrotron emission was not able toexplain the high optical ﬂux.
Except for the repetitive variations, stochastic variations were detected. We identiﬁed the two intervals
in which the random optical and X-ray variability werecorrelated. As a result of Spearman’s rank tests, the correlations were proved to be statistically signiﬁcant. We
estimated time delays between our optical light curves
and the INTEGRAL IBIS/ISGRI X-ray light curves withtime bin size of 5 s in the 25–60 keV energy band for thetwo terms (1) BJD 2457388.18–2457388.22 and (2) BJD
2457388.246–2457388.292 using a Bayesian method thatwas originally proposed to estimate time delays between
gravitationally lensed stochastic light curves (Tak et al.
2016). The 5-s binned X-ray light curves were derived
with the tool ii light. This method assumes that theirregularly sampled light curves are generated by a latent continuous-time damped random walk (DRW) process (Kelly et al. 2009) and that one of the latent light
curves is a shifted version of the other by the time lagin the horizontal axis and by the magnitude oﬀset in thevertical axis (Pelt et al. 1994). The model also adopts
heteroskedastic Gaussian measurement errors. We used

Table 1. Bayesian estimates of the time delays for term (1) BJD
2457388.18–2457388.22 and (2) BJD 2457388.246–2457388.292.
(We report posterior medians because posterior means are notreliable indicators for the center of a multi-modal distribution.
The posterior mode of time delays and median are identical up
to three decimal places for term (1). For term (2), the posteriormode is −33.2 s. ) The 68% interval indicates the quantile-based
interval and the 68% HPD interval represents the highest posterior
density interval.

Terms
(1)
(2)

Median−45.4 s−33.1 s68% Interval(−46.5 s, −45.1 s)(−33.3 s, −32.8 s)

68% HPD Interval(−45.6 s, −45.0 s)(−33.4 s, −32.9 s)

an R package, timedelay, which is publicly available at
CRAN2 , to implement the Bayesian model via a Markov
chain Monte Carlo (MCMC) method. Before the estimations, we examined the correlations between the opticaland X-ray luminosity with power law regressions andscaled the X-ray light curves using the estimated coeﬃcients of the regressions to meet the assumption in Tak
et al. (2016). The estimation results are summarized inTable 1. We obtained the ∼30–50 s X-ray delay to theoptical emission. The Gelman-Rubin convergence diagnostic statistics (Gelman & Rubin 1994) were 1.0004 and1.0009 in terms (1) and (2), respectively, close enough tounity.
4. Discussion
4.1. Origin of Optical Variations with Regular Patterns
The repetitive variations, especially the dip-type oscillations, are observed in both June and December outbursts
in 2015. First, the X-ray spectral analyses of the simultaneous X-ray data in the June outburst indicate that there*2 <https://cran.r-project.org/package=timedelay>

72

was no tendency for increased absorption when the X-ray
ﬂux decreased, suggesting that these dips do not originate from absorption (see Sec. 3 of Methods in Kimura etal. (2016)). In some epochs, we found evidence for heavy
obscuration as found in the GINGA data during the 1989
outburst (Życki et al. 1999); however this is not related
to dip-type variations. We can thus infer that the shortterm ﬂuctuations directly reﬂect variations in radiationfrom the accretion disc or its associated structures. Second, detailed analyses of the typical simultaneous broadband SED analyses in the June outburst show that the
majority of the optical ﬂux is most likely produced byreprocessing of X-ray irradiation in the disc. Third, we
detected the ∼1 min delay of optical lights to the Xray ones is consistent to the estimetd disc size. Thus, weconclude that the optical ﬂux was mainly originated fromthe X-ray reprocessing at the outer accretion disk when
the dip-type oscillations were observed in V404 Cyg.
In GRS 1915+105, it has been proposed that the observed X-ray repetitive variability (Belloni et al. 2000)
is caused by limit-cycle oscillations in the inner accretion disc due to the Lightman-Eardley viscous instability(see e.g., Janiuk & Czerny (2011)), which can explain aslow rise in brightness (mass accumulation) followed bya sudden drop (accretion to the black hole). Such models assume that the black hole is accreting mass nearlyat the Eddington rate. The theory is not adopted to therepetitive optical variations at low luminosity in V404
Cyg. We, therefore, suggest that an unknown instability
is related to the variability. We propose that they arecaused by disc properties which are inherent in blackhole binaries with long orbital periods.
4.2. Origin of Short X-ray Delays
To explain the X-ray delay of ∼30–50 s detected in the
December outburst to the optical emission, we consider
the condition that the disc was composed of an advection dominated accretion ﬂow (ADAF) and a thin disc.
We assume that the thin standard disc was extended tothe transition radius Rtr and that there was the ADAF
in the inside of the radius. In the ADAF, the mattermoves to the central object with the free-fall time scale(Narayan & Yi 1994). The speed in which the matter
falls approximately follows the equation, vf ∼ αvﬀ . Here,vﬀ and α represent the free fall speed and the viscousparameter. If we consider that the optical ﬂare was triggered at the thin disc and that the hot region accreted
to the central black hole with the free-fall time scale,the optical emission source can be located at the radiuslarger than 6.5 × 108 [cm] and the value of Rtr was estimated to be ∼6.5–8.5×108 [cm] for α = 0.1. This value
corresponds to ∼300rs . Here, rs (= 2GM/c2 ) represents
the Schwarzschild radius for a 9M black hole. Whenthe bolometric luminosity and the coeﬃcient of radiation eﬃciency (η) are ∼0.1LEdd and ∼0.1, respectively,

the temperature at the radius is estimated to be about
104 [K] which is consistent to the temperature in opticalemission region in a standard disc. The spectral statein this picture is expected to be hard and the low/hard
spectral state during 8.15–10.32 days (Motta et al. 2016)
supports this picture. In addition, the estimated valueof Rtr is as much as the value of the inner disc radius estimated with the SED analyses in the June outburst (seeSec. 8 of Methods in Kimura et al. (2016)). Moreover,the smaller amplitude of the optical variations than thatof the X-ray ones can be explained by the presence ofthe optical continuum emission at the outer disc.
We acknowledge the variable star observations from theAAVSO International Database contributed by observers
worldwide and used in this research. We also thankthe INTEGRAL groups for making the products of the
ToO data public online at the INTEGRAL Science Data
Centre. This work was supported by a Grant-in-Aid
from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan (25120007 TK)
and (26400228 YU). We are thankful to many amateur
observers for providing a lot of data used in this research.
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73

Two optical emission components with diﬀerent variability in V404 Cygni
Y. Tachibana1 , T. Yoshii1 , H. Hanayama2 , and N. Kawai1
1

2

Department of Physics, Tokyo Institute of Technology, 2-12-1, Ookayama, Tokyo, Japan
Ishigakijima Astronomical Observatory, National Astronomical Observatory of Japan, Ishigaki, Okinawa, Japan
E-mail(YT): tachibana@hp.phys.titech.ac.jp
Abstract
In this paper, we report the multi-color optical observation of V404 Cygni at the beginning of its
2015 outburst performed by the MITSuME and MURIKABUSHI telescopes. With time domain analysisof the multi-color light curves, we successfully decomposed optical variations into two components: ahighly-variable, and a less-variable component. The loci of the less-variable component in the color-color
diagram are consistent with that of multi-color blackbody radiation, while those of the highly-variable
component trace out those of power-law spectra with spectral indices α ∼ 0.6–1.0. For the less-variable
component, we argue that an irradiated disk with the innermost temperature higher than ∼ 2.0 × 104
K and the outermost temperature lower than ∼ 6.5 × 103 K is the most plausible source. The observed
spectral energy distribution (SED) from the optical to ultraviolet can be expressed by a model consistingof a power-law component and an irradiated disk component.
Key words: Black hole physics — X-rays: binaries — X-rays: individual (V404 Cygni)

1. Introduction
On 2015 June 15, V404 Cygni started an outburst after
26 years of quiescence. The burst was detected and reported by Swift/BAT and MAXI (Barthelmy et al. 2015,
Negoro et al. 2015).
In contrast to typical black hole transients, which exhibit a fast rise and exponential decay proﬁle during an
outburst, V404 Cygni showed multiple ﬂaring activity in
optical to γ-ray band during the outburst (Tanaka andShibazaki. 1996, Ferrigno et al. 2015). The nature of
the optical variation is still under discussion.
Kimura et al. (2016) argued that the source of the
optical emission is mainly the outer disk irradiated by
X-ray emission generated in the inner disk. This idea isalso supported by Muñoz-Darias et al. (2016). An altanative scenario that relativistic non-thermal emission(jet) component combined with the accretion disk emission contributes to the optical light, is proposed by Martiet al. (2016), Tanaka et al. (2016), Gandhi et al. (2016),
Lipunov et al. (2016), and Shahbaz et al. (2016).
In this paper, we present the multi-color optical light
curves obtained using with the MITSuME instrument
and results of the spectral-timing analysis.
2. Light Curve of V404 Cyg
Fig.1 shows the optical (g , RC , and IC -band) light curves
during MJD 57192–57194, corresponding to the 4–6 days
after the detection of X-ray outburst on MJD 57188.

Variation on MJD 57192 seems to be composed of two
distinct ﬂuctuations; big-slow swings and small-fast wiggles. On the other hand, in MJD 57193 and 57194, fast
drops in ﬂax rather than big-slow swings were seen, whilesmall wiggles remained. In all these variations, the ﬂuxes
in the optical three bands were always well correlated.
3. Analysis and results
3.1. Color-color Diagram
A color-color diagram is helpful to categorize the observed ﬂux variation and to test the applicability of assumed emission models. We therefore calculated thespectral indices between g -band and RC -band (αgr ),and RC -band and IC -band (αri ) for all the data point
in Fig.1 and presented them in a color-color diagram(Fig.2). When converting the observed magnitudes tointrinsic ﬂuxes Fν , we adopted AV = 3.3 and RV =AV /E(B−V)=2.8. The result is drawn by grey points
in Fig.2. The grey dashed line and black dashed line
trace loci for power-law spectra with diﬀerent spectral
indices and single temperature black body with diﬀerent temperatures, respectively, whereas the black line
and black-edged white line represents those for standard
disk blackbody and irradiated disk blackbody (p=3/4:
Shakura & Sunyaev 1973, and p=1/2: King & Ritter
1998) for diﬀerent inner/outer-most temperatures under
the condition that the disk temperature cannot acrossthe critical temperature ∼ 5000 K where the disk would

74

A

8

B

C

2.5

D

10

2.0
Ic
Rc

14

Magnitude

16

g'

2.55

2.60
E

8

2.65

2.70

F

2.75

2.80

G

12
14
3.55

3.60

3.65

H

8

3.70

3.75

3.80

Ic
Rc

−0.5

g'

−1.0

3.85

Ic
Rc

12

g'

14
4.55

4.60

4.65

0.5

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Single temperature black bodyMultiple tempetature blabk body (p=3/4)
Multiple tempetature blabk body (p=1/2)
Power−law−2.0 −1.5 −1.0 −0.5

0.0

0.5

1.0

1.5

α[g', Rc]

J

I

10

16

1.0

0.0

10

16

2.85

1.5

α[Rc, Ic]

12

4.70

4.75

4.80

4.85

MJD  57190 (days)
Fig. 1. The observed g , RC , and IC light curves of V404 Cyg in
MJD 57192–57194.

be unstable (King & Retter 1998, Życki et al. 1999),
respectively. Here, p is the gradient index of the surface
temperature of the accretion disk (T ) when expressed asa function of the radial distance from the center (r); i.e.
T ∝ r−p (e.g. Mineshige et al. 1994).
Apparently, the points in the color-color diagram does
not trace any of the model lines. It indicates thatthe optical variation cannot be attributed to a single
emission component that can be expressed by a singletemperature or multi-temperature blackbody, or powerlaw model.
3.2. Flux-ﬂux PlotUsing the color-color diagram, we showed that the optical variation of V404 Cyg cannot be attributed tochanges of a parameter of a single emission mechanism.
In this section, we further investigate the properties ofthe optical behavior phenomenologically on the ﬂux-ﬂux
plot.
Examples of ﬂux-ﬂux plots are shown in Fig.3 withlinear and broken linear model ﬁts for the same data
segments in the upper and lower panels respectively.
The three columns correspond to the episode containing
drastic ﬂux declines: MJD = 57192.70 – 57192.74 (B),57193.61 – 57193.64 (F), and 57194.66 – 57194.67 (I).
We adopt a broken-linear model to express the behaviorof the data on the ﬂux-ﬂux plot, since the broken-liner

Fig. 2. Color-color diagram for the data in Fig.1 with the spectral
indices between g -band and RC -band, and RC -band and IC -band.

model ﬁts gives signiﬁcantly small χ2 values. Two linear
branches below and above the break imply the existenceof two variable components with distinct spectral indices
α because the slope of the linear model on the ﬂux-ﬂux
plot corresponds to (ν2 /ν1 )α .
To derive slopes for the two branches, we ﬁtted a
broken-linear model to the data in entire periods A–
J. We found in ﬂux-ﬂux plots (e.g. Fig.3) that the dynamic range of the ﬂux variation in the component above
the break is lager than that below. It indicates thatone variable component has quite large variability anddominates the optical variation when the component ismuch brighter than another one. Therefore, we named
the components dominating above the break as “highlyvariable component” (HVC) and the one below as “lessvariable component” (LVC).
3.3. Decomposed Color-color Diagram
Fig.4 shows the color-color diagram for the decomposedcomponents; HVC and LVC. The spectral indices of the
HVC lie on the gray dashed line indicating power-law
spectra with variable indices. On the other hand, thoseof the LVC are consistent with emission from the standard disks (p=3/4) with innermost temperatures (Tin )of typically ∼ 1.5 × 104 K and 5000 K at the outer edge
(Tout ), or irradiated disks (p=1/2) with Tin  2 × 104 K
and Tout  6500 K.
4. Discussion
4.1. Optical Emission Components
We have shown that the optical variation is decomposed
to two components; less-variable component (LVC) andhighly-variable component (HVC). HVC has larger amplitude (∼5–10 times larger than that of LVC) and bluer

Normalized flux density (Ic)

75







4





MJD = 57192.70  57192.74 (B)Linear model

MJD = 57193.62  57193.64 (F)Linear model

MJD = 57194.66  57194.67 (I)Linear model

MJD = 57192.70  57192.74 (B)Broken linear model

MJD = 57193.62  57193.64 (F)Broken linear model

MJD = 57194.66  57194.67 (I)Broken linear model

3
2
1
0
10
0










Normalized flux density (Ic)

10
4
3
2
1
0
10
0

10
0

1
2
3
4
5
Normalized flux density (g')

0

1
2
3
4
5
Normalized flux density (g')

0

1
2
3
4
5
Normalized flux density (g')

Fig. 3. The examples of ﬂux-ﬂux plot with the ﬂuxes between g -band and IC -band. The upper panels are those ﬁtted with a liner model,
while the lower panels are ﬁtted with a broken liner model.

color, and thus corresponds to “big-slow swings”. On
the other hand, LVC which has a redder spectrum corresponds to a characteristic behaviors mentioned as “smallfast wiggles”. The latter is similar to “heartbeat-type”
oscillations observed by Kimura et al. (2016), considering its small amplitude and its dominance at the low ﬂux
phase.
4.2. The emission origin of LVC
We found that spectral indices of LVC are consistent withthose of a standard disk or an irradiated disk. However,a standard disk truncated at such a low innermost temperature (∼ 1.5×104 K) seems to be diﬃcult to reconcile
with the detection of a disk component in the soft X-ray
band (Radhika et al. 2016 and Walton et al. 2016). For
the origin of LVC, thus, the irradiated disk is favored.
From the observed ﬂux, the outer radius of the disk isevaluated to be ∼ 3–5 ×105 Schwarzschild radius. It isconsistent with the size limit based on the timescale of
LVC. The irradiated disk is therefore the most plausible
source of the LVC.
4.3. The emission origin of HVC
The decomposed color-color diagram shows that the
spectra of HVC follow power-law models with spectral

indices of ∼ 0.6–1.0.
Since the break of a jet spectrum between the optically thick and thin regions is expected to lie at ∼ 1.4–
4.7 ×1014 Hz (Shahbaz et al. 2016, Tanaka et al. 2016,
and Itoh et al. 2016), the blue spectrum seen in HVC(α ∼ 0.5–1.0) is possibly attributed to the transitioning spectrum from optically thick synchrotron spectrum(α = 2.5) to thin one. The jet is therefore one of thecandidate of the origin of HVC.
Another possibility is the cyclo-synchrotron emission
as proposed for XTE J1118+480 (Merloni et al. 2000
and Kanbach et al. 2001) and V4641 Sgr (Uemura
et al. 2002). The self absorbed regime of the cyclosynchrotron spectrum also can produce the hard spectral index than ∼0.6. The cyclo-synchroron emission is
also one of the possible interpretation of the power-law
component (HVC).
4.4. SED from optical to UV
To conﬁrm our interpretation, we tried to describe thespectral energy distribution (SED) from optical to ultraviolet with a power-law and a irradiated disk spectrum. Constructed SED at MJD 57194.616 (included in“H” portion) are displayed in Fig.5. Four S wift/UVOT
data (V, U, UVW1, and UVW2) are obtained quasi-

76

2.5

−22

1.5

α[Rc, Ic]

Log[ Fν (erg s−1 cm−2 Hz−1 )]

2.0

−21
Little variable component
Highly variable component

−23

1.0
0.5

−24

0.0

−25

−0.5

Single temperature black bodyMultiple tempetature blabk body (p=3/4)
Multiple tempetature blabk body (p=1/2)
Power−law−1.0−2.0 −1.5 −1.0 −0.5

0.0

0.5

1.0

−26

1.5

α[g', Rc]

Multiple temperature black body (p=1/2)
Power−law

−27
13

14

15

16

Log[ ν (Hz)]Fig. 4. Color-color diagram with the spectral indices between g -band
and RC -band, and RC -band and IC -band, after the decomposition.

simultaneously. The solid line and the dashed line denotethe irradiated disk spectrum (temperature Tout = 5×103
K with the size of the emission region R = 4 × 105 Rs )and power-law spectrum with the index α = 1.0 respectively. Although we did not use the UVOT data for
determination of the spectral model, all data points including the UV excess appear to be explainable by thetwo component model.
5. Summary
In this work, we found that the optical emission of V404
Cygni in the 2015 outburst can be decomposed to threecomponents: the less-variable component (LVC), thehighly-variable component (HVC), and the stable component, based on the ﬂux-ﬂux plots which is well represented by broken-liner expressions.
The decomposed spectral indices of LVC and HVCon the color-color diagram are consistent with those ofmulti-temperature black body spectra whose local temperature T (r) is proportional to r−1/2 and a power-law
spectra, respectively.
The inner- and outer-most temperatures (Tin/out ) andthe outermost radius of the irradiated disk are Tin 
2.0 × 104 K, Tout  6.5 × 103 K, and (3–5) ×104
Schwarzschild radius respectively. The index of thepower-law component is ∼0.6–1. We proposed a nonthermal jet or a cyclo-synchrotron emission as the originof the power-law component (HVC). The validity of ourinterpretation is supported by the spectral energy distribution from optical to UV with an irradiated disk and apower-law spectrum.
Our interpretation that the optical variation consistsof two variable components, presents a possible solution

Fig. 5. SED from optical to UV with an irradiated disk spectrum (solidline) and power-law spectrum (dashed line) in MJD ∼ 57194.61– 57194.62 (included in “H” period in Fig. 1).

for the problem pointed out by previous papers. Kimura
et al. (2016) presented a problem about the amplitudesof the variations in X-ray and optical which are generally
correlated; the X-ray ﬂux variations were much larger
than the optical ones. The smaller amplitude in the
optical-band compared to X-ray is naturally explainedby the “bias” (LVC and/or the stable component) underlying the component (HVC) correlated to X-ray variation.
References
Barthelmy S. D. et al. 2015, ATel, 17929, 1
Ferrigno C. et al. 2015, ATel, 7662, 1
Gandhi P. et al., 2016, MNRAS, 459, 554
Itoh R. et al. 2016, ArXiv, e-prints 1612.07513
Kanbach G., et al., 2001, Nature, 414
Kimura M. et al. 2016, Nature, 529, 54
King A. R. and Ritter, H. 1998, MNRAS, 293, L42
Lipunov V. M. et al., 2016, ArXiv, e-prints 1608.02764
Marti J. et al. 2016, A&A, 281, 1094
Merloni A., et al., 2000, MNRAS, 318, L15
Mineshige S. et al. 1994, ApJ, 426, 308
Muñoz-Darias et al. 2016, Nature, 534, 75
Negoro H. et al. 2016, ATel, 7646, 1
Radhika D. et al., 2016, ArXiv, e-prints 1601.03234
Shahbaz T. et al., 2016, MNRAS, 463, 1822
Shakura N. I., Sunyaev R. A., 1973, A&A, 24, 337
Tanaka Y. & Shibazaki, N., 1996, X-ray novae. Annu.
Rev. Astron. Astrophys. 34, 607
Tanaka Y. T. et al., 2016, ApJ, 823, 35
Uemura M. et al., 2002, PASJ, 54, L79
Walton D. J. et al. 2016, ArXiv, e-prints 1609.01293
Życki P. T. et al., 1999, MNRAS, 309, 561

77

Observation of the long-term variability of Cygnus X-1
with MAXI
Juri Sugimoto,1,2 Shunji Kitamoto,2 Tatehiro Mihara1 and the MAXI team
1

MAXI team, Institute of Physical and Chemical Research (RIKEN), 2-1 Hirosawa, Wako, Saitama 351-0198, Japan2
Department of Physics, Rikkyo University, 3-34-1 Nishi-Ikebukuro, Toshima, Tokyo 171-8501, Japan
E-mail(JS): sugimoto@crab.riken.jp
A BSTRACTWe studied the long-term X-ray variability of the black hole binary, Cygnus X-1, using about
6 years of MAXI data from 2009 to 2016, in the high/soft state (HSS) and the low/hard state(LHS). We analyzed an aperiodic long-term variation, using Normalized Power Spectrum Densities (NPSDs), in the frequency range from 10−7 to 10−4 Hz, with 2 − 4 keV, 4 − 10 keV and
10 − 20 keV data. From the comparison with previous works in the frequency region above
10−3 Hz, it was shown that the NPSD extends with a single index from 10−3 to 10−7 Hz in boththe LHS and HSS. Such the long-term variation is mostly caused by the power-law component
in the energy spectrum. In the LHS, the variability slightly decreased towards higher energies,
while in the HSS, the variability was found to increase signiﬁcantly with energy. To explain our
results of the long-term aperiodic variation, we propose a new large-scale disk structure. In ourmodel, an optically-thin and geometrically-thick accretion ﬂow coexists with an optically-thickand geometrically-thin, standard accretion disk, which is extending up to a large radius of ∼ 1012
cm, in both HSS and LHS.
K EY WORDS : black hole physics — accretion, accretion disks—X-ray:general — stars: individual: Cygnus X-1

1. Introduction
In the X-ray universe, a mass accretion is seen invarious phenomena. Among them an accretion disk
is commonly formed around a compact star in an
X-ray binary system. Such a compact star shines releasing the gravitational energy of the accreting gas
from the companion star. The X-ray binary, especially with a black hole (BH), is studied by a lot of
researchers, focusing on the accretion disk, including the accretion ﬂow.
An important information obtained from observations is a time variation of the X-ray intensity. In
fact, the typical BH binary, Cygnus X-1 (Cyg X1), is drawn attention to its violent time-variation
since its discovery in 1960’s. It has an irregular andlarge-amplitude variation in the range of 0.1 ∼ a
few second. The cause of the time variation is related with the characteristics of the accretion ﬂow in
the accretion disk, especially within 10 rS (rS is the
Schwarzschild radius) where majority of X-rays are

emitted.
On the other hand, the long-term variation morethan 1000 s is not well understood. Because most ofobservations are made by X-ray telescopes mounted
on the satellites orbiting a low Earth orbit. Such atime variation in the scale of several tens of days
is important to understand the complete structure of
the accretion disk. In order to investigate the overall structure of the accretion disk, especially of Cyg
X-1, we studied aperiodic long-term variation of theintensity. We use the X-ray data obtained by MAXI
which provides us a statistically good X-ray intensity and a spectrum every 90-minite spaning morethan six years.
2. Observation
We analyzed the data of Cyg X-1 obtained withMAXI (Matsuoka et al. 2009). We downloaded
archival data of Cyg X-1 of one-day bin and 90-

78

minute bin from the MAXI home page1 .
Figure 1 showed one-day bin light curves of Cyg
X-1 obtained with the GSC from 2009 August 15(MJD = 55058) to 2016 January 30 (MJD = 57417),
in three energy bands (2 − 4 keV, 4 − 10 keV and
10 − 20 keV). A time history of the hardness ratio(HR), I(4 − 10 keV) / I(2 − 4 keV), is also plotted.
The state of Cyg X-1 can be recognized by the values of the HR as the HSS (HR<0.43) and the LHS(HR>0.48).
Since the start of the MAXI observation in 2009
August, Cyg X-1 was in the LHS for about ten
months. At around MJD = 55378, the source madea transition to the HSS, where it stayed for anotherten months. After several times of state transitions,
it has been mainly in the HSS since MJD = 56078.
3. Analysis
3.1. Deﬁnition of NPSDs
The X-ray light curve with 90-minute bin was converted to the PSD by the discrete Fourier transformation as,

Fc (f ) =Fs (f ) =2
yj sin(2πf Δt × j),
N j=1

N

(1)

where yi is the intensity (photons s−1 cm−2 ) at thei-th bin, N is the number of data after ﬁlling thegaps as described in Sugimoto et al. (2016), andΔt is a time interval of 5400 s, which is the regular sampling time of the MAXI public data, f is
the frequency which is an integer multiple of Δf = T1
and T is the total time span of the observation. The
factor of 2/N is employed so that Fc (f ) and Fs (f )represent the amplitude of the cosine and sine components, respectively, regardless of N . The power,
12
2
2 (Fc (f ) + Fs (f )), is that in the frequency range
Δf . Therefore the PSD, power per Hz, is,P (f ) =1F 2 (f ) + Fs2 (f ) .
2Δf c

To obtain the real NPSD from the MAXI data
two corrections are necessary; the gap correction
and the aliasing correction. The MAXI data are notcompletely regular-sampled. Sometimes data gaps
are caused by sun avoidance, high particle background regions, small dead regions around the scanpoles which move with the precession period of theorbit of the ISS, and other effects. Furthermore,
MAXI measurements are snapshots. It scans an Xray source only for 40 ∼ 70 s, every 5400 s period. Each data point is a very short snapshot, witha long interval to the next sampling. As a result, a
MAXI data point contains source variations not only
on time scale longer than 5400 s, but also to these inbetween 40 ∼ 70 s and ∼ 5400 s. We described theestimation of these effects and derived a “transfer
function” for the calculated PSD to the best estimation of the real PSD in Sugimoto et al. (2016).

3.2. Long-term PSDs v.s. short-term PSDs

N

2
yj cos(2πf Δt × j),
N j=1

average intensity following Miyamoto et al. (1994).
It is called an Normalized PSD (NPSD).

(2)

The unit is (RMS2 /Hz). To represent relative amplitude, the PSD is normalized by the square of the*1 http://maxi.riken.jp/top/index.php?cid=
1&jname=J1958+352

Although we can make a long-term PSD by using
all the 6years of data, single PSD is difﬁcult to estimate the errors of data points. We shorten the data
length to 40 and 43 days for the LHS and the HSS,
respectively, calculated NPSDs, and averaged them.
The errors of NPSDs obtained in this way can be expected statistically and have smaller errors. Instead,it lacks the information in the lowest frequencies.
The data span covers a frequency range down to ∼3×10−7 Hz. Six data segments were extracted for
the LHS, and eight data segments for the HSS. Then,
we converted these data segments individually into
NPSDs, and took their averages over the LHS andHSS separately.
Figure 2 shows our 4 − 10 keV NPSDs in bothstates, in comparison with previous works with
RXTE data by Cui et al. (1997), Pottschmidt et al.
(2003), and Reig et al. (2002). In both states, our results at 10−4 Hz locate on simple extrapolations of
the NPSDs obtained previously by the RXTE/PCA
in the frequency region above 10−3 Hz. The NPSD
in Reig et al. (2002), which is obtained by the
RXTE/ASM, is located between the two states ofthe present work. This difference is presumably because their data included some of the HSS and thetransition periods, although the main part was in the
LHS.

79

2-4 keV

4

X-ray intensity(photons s-1 cm-2 )

2
0
1.5
1
0.5
0

4-10 keV

10-20 keV

0.4
0.2
0

(4-10 keV)/(2-4 keV)

1
0.5
0

55200

2010

55400

55600

2011

55800

56000

56200

2012

56400

2013

56600

56800

2014

57000

57400 (MJD)

57200

2015

2016 (year)

Fig. 1. One-day bin light curves and HR histories of Cyg X-1 obtained with the MAXI/GSC. From the top to bottom panels,
the 2 − 4 keV, 4 − 10 keV and 10 − 20 keV intensities, and the hardness ratio, I(4 − 10 keV) / I(2 − 4 keV), are plotted.
The blue and red regions indicate the LHS and the HSS periods, respectively. Data points with large error bars, due
to high background counts, were omitted.

3.3.

Energy dependence of the long-term
NPSD in the HSS
The long-term NPSD in the HSS was obtained forthe ﬁrst time. Figure 3 shows the NPSD in the
HSS in the frequency range of 3 hours to 1 month.
The vertical error bars represent the NPSD scattering (standard deviation) within each ensemble. We
found that the high-energy bands (4 − 10 keV and
10 − 20 keV), which correspond to the power law
component in the energy spectrum, are more variable than the low energy band (2 − 4 keV), whichcorresponds to the disk component. It means thatthe power law component is responsible for the observed long-term intensity variations, while the disk
emission is essentially constant.
4. Discussion
The short-term variation above 10−3 Hz is thought
to reﬂect the phenomena occurring in the vicinityof the BH. On the other hand, it is unlikely that thelong-term variation in the range of 10−7 Hz is generated there. We consider that it is generated at alarge radius and propagates inwards as a variationof the mass accretion rate. Then most of the X-rays
are radiated from the vicinity of the BH. Here, we
discuss that how the long-term variation is propagated to the vicinity of the BH. The time scales ofthe variation newly produced in the accretion disk
are the dynamical time scale, td , and thermal timescale, tth , expressed as
td = Ω−1K ,tth = (αΩK )−1 .

(3)

On the other hand, the time scale of the accretionis the viscous time scale of the accretion disk, tvis ,expressed astvis = (αΩK )

−1



H
r

−2
.

(4)

We compared these time scales in the standard
disk and in the RIAF. In an optically-thick andgeometrically-thin standard disk, the thickness ofthe disk H is much smaller than the radius r. Then,the viscous time scale is larger than the other twotime scales, and variations produced at any outerradii would be strongly dissipated and would notpropagate down to the X-ray emitting region. On theother hand, in an optically-thin and geometricallythick ﬂow (RIAF), H is the same level as r. In thiscase the thermal time scale is of the same order as
the viscous time scale, and the ﬂuctuations can propagate inwards, varying the X-ray intensity on a wide
range of time scales.
Then, we proposed a new large-scale accretion disk structure as illustrated in ﬁgure 4. The
optically-thin ﬂow (RIAF) should co-exist from thevicinity of the BH to the radius of ∼ 1012 cm.

80

References
Churazov, E. et al. 2001, MNRAS., 321, 759
Cui, W. et al. 1997, Transparent Universe, 382, 209
Matsuoka, M. et al. 2009, PASJ, 61, 999
Mihara, T. et al. 2011, PASJ, 63, 623
Miyamoto, S. et al. 1994, ApJ, 435, 398
Pottschmidt, K. et al. 2003, A&A, 407, 1039
Reig, P. et al. 2002, A&A, 383, 202
Sugimoto, J. et al. 2016, PASJ, 68, S17
Tomida, H. et al. 2011, PASJ, 63, 397

105

Reig+2002

103

HSS
LHS

101

Cui+1997

10-1
Pottschmidt+2003
10-3
10-8

10-6

100
10-4
10-2
frequency (Hz)102

Fig. 2. A comparison of the MAXI 4 − 10 keV NPSDs in the
LHS (blue) and the HSS (red), with the previous works.
The dashed line is the 2 − 13 keV NPSD in the LHS
by Pottschmidt et al. (2003), the dotted line is that in
the HSS in 2 − 13 keV by Cui et al. (1997), and the
dot-dashed line is that by Reig et al. (2002), where themost data are in the LHS.

105

NPSD

5. Conclusion
Using six years archival data obtained with MAXI,
we derived the NPSDs of Cyg X-1 in its LHS and
HSS from 10−7 Hz to 10−4 Hz in the three energy
bands. It is paticularly important that the long-term
variations in the two states were studied in a uniﬁed
way using the same instrument, with the same analysis method, and under similar data statistics. The
PSD of the both states in the frequency range between 10−5 and 10−4 Hz were also provided by ourwork for the ﬁrst time. Our results in both states
were approximated by a power law, and on an extrapolation of the NPSD previously obtained in frequencies above 10−4 Hz.
Also the long-term NPSD in the HSS was obtained for the ﬁrst time. The NPSD from 10−7 Hz
to 10−4 Hz in the HSS depends on energy in sucha way that it is 5∼6 times higher in the 10 − 20
keV band than that in the low energy band (2 − 4keV). This result is expected to provide some clues
to the still unknown origin of the hard-tail component, which is nearly always observed in the HSS.
We then proposed a new disk structure. The
optically-thin corona should extend to ∼ 1012 cm,
and a standard optically-thick and geometricallythin disk co-exists in the out side of the certain radius. The long-term variation produced at the large
radius is transported to the vicinity of the compact
star, via the optically-thin corona, and eventually Xray intensity of the power law component shows thelong-term variation.

NPSD

The long-term variation produced at the large radius is transported to the vicinity of the compact
star, via the RIAF, and eventually X-ray intensity ofthe power law component shows the long-term variation.

104
103
102
101
10

month
-7

week

day

-6

10
10-5
Frequency (Hz)10-4

Fig. 3. The NPSDs calculated over data segments (43 d ×
8 for the HSS), and averaged.

X-ray emitting region

RIAF
BH

standard disk~1012 cm

Fig. 4. Illustration of our accretion model.

81

Accretion Flow Properties of three MAXI Black Hole Candidates:
Analysis with the TCAF Solution
Dipak Debnath,1 Sandip K. Chakrabarti,2,1 Arghajit Jana,1 D. Chatterjee1 A. A. Molla,1 S. Mondal3,11Indian Center for Space Physics, 43 Chalantika, Garia St. Rd., Kolkata, 700084, India.
2
S. N. Bose National Centre for Basic Sciences, Salt Lake, Kolkata, 700098, India.
3
Instituto de Fı́sica y Astronomı́a, Universidad de Valparaı́so, Gran Bretana N 1111, Playa Ancha, Valparaı́so, Chile.
E-mail: dipak@csp.res.in (DD); chakraba@bose.res.in (SKC); argha@csp.res.in (AJ)
Abstract
Recently after the implementation of Chakrabarti and his collaborators Two-component advective ﬂow(TCAF) model into HEASARC’s spectral analysis software package XSPEC as an additive table model,
we found that it is quite capable to describe the underlying accretion ﬂow dynamics around black holes
with spectral ﬁtted physical parameters. Properties of diﬀerent spectral states and their transitions during
an outburst of a transient black hole candidate (BHC) are more clearly understandable. One can evenpredict frequency of the dominating quasi-periodic oscillations (QPOs) from the TCAF model ﬁtted shock
parameters and even predict the most probable mass range of an unknown BHC from TCAF ﬁts. This
gives us a conﬁdence that the description of accretion process is more clear than ever before. Recently
we studied three MAXI/GSC discovered Galactic transient BHCs: MAXI J1659-152, MAXI J1836-194,
MAXI J1543-564, with the TCAF solution to study accretion ﬂow properties of these compact objects
during their very ﬁrst outbursts. We classiﬁed entire outbursts of these sources into diﬀerent spectral
states on the basis of nature of variations of TCAF model ﬁtted physical ﬂow parameters and observed
QPOs. Probable mass ranges of these BHCs are also estimated from our study. Estimation of X-ray jet
ﬂuxes during the 2011 outburst of MAXI J1836-194 are also done.
Key words: X-Rays:binaries – stars individual: (MAXI J1659-152, MAXI J1836-194, MAXI J1543564) – stars:black holes – accretion, accretion disks – ISM: jets and outﬂows – radiation:dynamics

1.

Introduction

Transient black hole candidates (TBHCs) are very interesting objects to study in X-rays as they exhibit
rapid evolutions in their temporal and spectral properties during outbursts, which are strongly correlatedwith each other. Several works are present in the literature (see, e.g., Debnath et al., 2008, 2013 and references therein), which explain variations of spectral
and temporal properties of these objects during their
X-ray outbursts. In general, it has been found thatthese objects exhibit diﬀerent spectral states, such as
hard (HS), hard-intermediate (HIMS), soft-intermediate(SIMS) and soft (SS) during their outbursts. They also
exhibit low and high frequency quasi-periodic oscillations(QPOs) in some of these spectral states (see, Debnath
et al. 2013 and references therein). The frequencies of
QPOs are found to evolve with time during the rising anddeclining phases (more precisely, during hard and hardintermediate spectral states) of the outbursts. According to the nature of the variations of temporal and spec-

tral properties during outbursts, we can classify transient
BHCs into two types: i) type-I or classical-type, where
all spectral states are observed (HSSS via two intermediate, HIMS and SIMS spectral states), and ii) type-IIor harder-type, where SS (sometimes even SIMS) could
not been observed. The type-II type of outbursts are
generally termed as ‘failed-outbursts’.
It is well known that emitted spectrum of the radiation coming from BHCs is a multi-color in nature, andcontains both thermal and non-thermal components.
One component is basically the multi-color blackbody
radiation from the standard Keplerian disk (Shakura& Sunyaev 1973), and the other is a power-law component, which is originated from the ‘Compton’ cloud(Sunyaev & Titarchuk 1980) composed of hot electrons,
which is cooled down due to repeated Compton scattering of softer (lower energy) photons. In the two component advective ﬂow (TCAF) solution of Chakrabarti& Titarchuk (1995; hereafter CT95), the role of socalled ‘Compton cloud’ or ‘hot corona’ is replaced by the82

CENtrifugal pressure supported BOundary Layer (CENBOL), which automatically forms behind the centrifugal barrier due to piling up of the low viscosity (subcritical) and low angular momentum optically thin matter, known as sub-Keplerian (halo) accretion component.
In TCAF, the other component of accretion ﬂow is thehigh viscous, optically thick and geometrically thin Keplerian (disk) component which is submerged inside thehalo component. CENBOL being much hotter in general, the Keplerian disk is naturally truncated at theshock location close to the black hole. This Keplerian
ﬂow settles down to a standard SS73 disk when coolingis eﬃcient. Low energy (soft) thermal photons from the
Keplerian disk intercept with the CENBOL (composedof hot electrons) and emit as high energy (hard) photons
by cooling down CENBOL through repeated Compton
scatterings.
Recently, in 2014, Debnath and his collaborators
have successfully been implemented TCAF solution into
HEASARC’s spectral analysis software package XSPEC
as an additive table model to ﬁt black hole spectrum.
After that, they have been quite successful to explain accretion ﬂow properties of few transient BHCs (such as,
GX 339-4, H 1743-322, Swift J1753.5-1027, other than
three MAXI objects, MAXI J1659-152, MAXI J1836194, MAXI J1543-564) during their X-ray outbursts(Debnath et al., 2014; 2015a,b, 2017; Mondal et al.,2014; Jana et al., 2016; Chatterjee et al. 2016; Molla
et al. 2016, 2017; Bhattacharjee et al. 2017). From
the TCAF model ﬁtted spectrum, one can not only directly obtain information about the instantaneous accretion rates of the two components (Keplerian disk rate ṁdin ṀEdd and sub-Keplerian halo rate ṁh in ṀEdd ), butalso obtain crucial information on shock parameters (instantaneous shock location Xs in Schwarzschild radiusrg =2GM/c2 , and shock compression ratio R), which allows us to estimate frequencies of the dominating QPOs
(if observed in PDS; see, Debnath et al. 2014). Depending upon the variations of accretion rate ratio (ARR =halo/disk rates) and nature of QPOs (shape, frequency,
Q value, rms%), classiﬁcation between diﬀerent observed
spectral states and their intermediate transitions are wellunderstood. A strong correlation between the spectral
and timing properties during the outburst period in theform of hysteresis loop, i.e, in ARR-intensity diagram(ARRID) was observed, where diﬀerent spectral states
are found to be correlated with diﬀerent branches of the
diagram (see, Jana et al. 2016). Most probable mass
(MBH ) ranges of MAXI J1659-152, MAXI J1836-194,
MAXI J1543-564, H 1743-322, etc. have been predictedfrom our spectral analysis with TCAF model ﬁtted constant normalization (N ) method (see, Molla et al. 2016,2017; Jana et al. 2016; Chatterjee et al. 2016; Bhattacharjee et al. 2017).

In this Conference Proceeding, we will brieﬂy discuss
about the accretion ﬂow properties of three MAXI objects (MAXI J1659-152, MAXI J1836-194, MAXI J1543564) during their very ﬁrst outbursts after discovery, obtained from our spectral and temporal analysis under
TCAF paradigm (Debnath et al. 2015; Molla et al. 2016;
Jana et al. 2016, 2017; Chatterjee et al.2016).
MAXI J1659-152: (Debnath et al. 2015b; Molla etal. 2016)
The source was discovered by MAXI/GSC on 2010 Sep.
25 at R.A.= 16h 59m 10s , Dec = −15◦ 16 05 (Negoro etal. 2010). It is shortest orbital period (2.414 ± 0.005 hr)
BH with acceptable ranges of distance, inclination angle,
mass are 5.3−8.6 kpc, 60◦ −80◦ , 3−8 M respectively. In
Debnath et al. (2015b), we studied both the spectral andtiming properties of the source during its 2010 outburst
by analyzing 2 − 25 keV RXTE/PCA data of 30 observational IDs, starting from 2010 Sep. 28 (MJD=55467)
to 2010 Nov. 11 (MJD=55508). 2.5 − 25 keV spectra are
ﬁtted with TCAF model ﬁts ﬁle by keeping all model
parameters (ṁd , ṁh , Xs , R, N ) as free, except MBH
was kept frozen at 6 M . Depending upon the variations of ARR and nature of QPOs, we classiﬁed entire
2010 outburst into three spectral states (HS, HIMS and
SIMS) in the sequence of: HIMS (Ris.)→SIMS→HIMS(Dec.)→HS (Dec.). On the ﬁrst PCA observation day,source was already in HIMS (Ris.). Monotonic evolutions of low frequency QPOs during HIMS (Ris.) and
HIMS-HS (Dec.) are observed and during the SIMS,QPOs are observed sporadically. No, signature of SS isobserved during the outburst, which may be due highsupply of low viscous halo matter both via accretion andwind, since it is a short orbital period BH.
In Molla et al. (2016), we reﬁtted all spectra by keeping TCAF model normalization (N ) in a ﬁxed value(which is the average N value of the ﬁtted normalizations, when all parameters were kept free) to estimate
probable mass range of the source as 4.7 − 7.8 M . We
also veriﬁed this mass range with the estimated mass
value (5.1−7.4 M ) obtained from declining phase QPO
frequency evolution, ﬁtted with propagating oscillatory
shock (POS) model (Chakrabarti et al. 2005, 2008; Debnath et al. 2010, 2013; Nandi et al. 2012). Finally, combining estimated mass ranges obtained from TCAF and
POS models, we obtained a most probable mass rangeof the source as 4.7 − 7.8 M or 6+1.8−1.3 M .
2.

3. MAXI J1836-194: (Jana et al. 2016)
This short orbital period (< 4.9 hr), low inclination angle (4◦ − 15◦ ) angle and highly rotating (a = 0.88 ± 0.03)
BH was discovered on 2011 Aug. 29 simultaneously by
SWIFT/BAT and MAXI/GSC (Negoro et al. 2011a) at
R.A. = 18h 35m 43s .43, Dec = −19◦ 19 12 .1. Cenko et83

al. (2011) reported companion as a high massive B[e]star. The probable mass and distance of object werereported as 4 − 12 M , 4 − 10 kpc respectively. The
source was highly active in radios. In Jana et al. (2016),
we studied spectral and timing properties of the source
using RXTE/PCA data of 35 observational IDs starting from 2011 Aug. 31 (MJD=55804) to 2011 Nov. 24(MJD = 55889). 2.5 − 25 keV spectra are ﬁtted with
TCAF model ﬁts ﬁle to extract physical ﬂow parameters.
Low frequency QPOs are observed sporadically throughout the outburst with a general trend of increasing nature of frequency during the rising phase and decreasing
nature of frequency during declining phase. Depending
upon the variations of TCAF model ﬁtted ﬂow parameters, ARRs and nature of QPOs, we classiﬁed entire
2011 outburst of the source into two harder states (HSand HIMS). These spectral states are observed in the sequence of: HS (Ris.)→HIMS (Ris.)→HIMS (Dec.)→HS(Dec.). We divided HIMS into two classes, since during
the HIMS (Ris.), ARR decreased due to decreasing nature of ṁh and increasing trend of ṁd . On the transition
day between two HIMS (MJD=55820), ARR becomesminimum (due to maximum value of ṁd , and minimum
value of ṁh ) with maximum value of the observed QPOs.
After that ARR increased and becomes maximum on thetransition day between HIMS (Dec.) to HS (Dec.) due torise in ṁh . There after during HS (Dec.), ARR decreased
due to fall of both rates. During entire 2011 outburst,
no signature of softer spectral states (SS and SIMS) are
observed, which may be due to high supply of low viscous wind matter, since BH accretion disk is immersed
inside the excretion disk of the B[e] companion.
To ﬁnd correlation between spectral and timing properties, we draw ARRID and compared it with hardnessintensity diagram (HID; Debnath et al. 2008 and references therein). In ARRID (see, Fig. 4 of Jana et al.
2016), we plotted 2 − 25 keV PCA count rates as a function of ARR. Diﬀerent observed spectral states and theirtransitions are found to be more prominent in ARRID
compared to HID (N.B.: point D, where transition between declining HIMS and HS is feasible only in ARRID). Observing peak diﬀerences between disk and halo
rates during the entire outburst, we predicted viscoustime scale as ∼ 10 days. From our detailed spectral
study with the TCAF solution, we also estimated probable mass range of the source as 7.5 − 11 M , since inthe model mass is an input parameter. In Jana et al.
(2017), we also calculated X-ray jet ﬂuxes from spectral
analysis using TCAF solution.

where, νs0 is the ﬁrst observed QPO frequency and tinf allis the infall time.
In Chatterjee et al. (2016), we also estimated probablemass range of the source from our spectral and temporal
analysis with the TCAF solution and the POS model
respectively. From TCAF model spectral ﬁts, MBH wasobserved in the range of ∼ 13.5−14.0 M , where as fromthe POS model ﬁtted QPO frequency evolution MBH
was found at 13 M for the best-ﬁt (χ2red =0.90). Thenwe varied MBH in POS model code for both positive andnegative values and observed, for 90% conﬁdence (χ2red ≤2.7) allowed mass range could be ∼ 12.6−13.6 M . Thenﬁnally combining results obtained from both TCAF and
POS models, we predicted mass range of the source as∼ 12.6 − 14.0 M or 13+1.0−0.4 M .

4. MAXI J1543-564: (Chatterjee et al. 2016)
The source was discovered by MAXI/GSC on 2011 May
8 at R.A.=15h 43m 9s .12 and Dec.=56◦ 25 15 .6 (Negoro
et al. 2011). Stiele et al. (2011) studied entire outburst

5. Discussions and Concluding Remarks
We studied spectral and temporal properties of three
MAXI observed BHCs during their very ﬁrst outbursts
after discovery using RXTE/PCA data. Out of themof the source using spectral ﬁts with combined diskbb
and simpl models, where as in Chatterjee et al. (2011),
we studied initial rising phase of the outburst with the
TCAF solution. We analyzed initial seven observationsof RXTE/PCA in the rising phase of the outburst, starting from 2011 May 10 (MJD=55691.09) to 2011 May
15 (MJD=55696.66). On the ﬁrst observation day, thesource was already in HIMS. From the nature of the variations of TCAF model ﬁtted physical ﬂow parameters,
ARRs and QPOs, during we classiﬁed our observation
period into two intermediate spectral states, HIMS and
SIMS. Low frequency QPOs are observed during ﬁrst six
observations. During initial ﬁve observations a monotonic evolution (increasing trend) of type-C QPOs (from1.05− 5.70 Hz) is observed and on sixth day, it decreased
to 5.08 Hz. We deﬁned ﬁfth observation as the transition
day between two outbursts. From TCAF model ﬁts, Keplerian disk rate (ṁd ) increased as outburst progressed,
where as sub-Keplerian halo rate (ṁh ) decreased initially upto the transition day and then slightly increased
for last two observations. Shock is found to move inward with decreasing strength. Evolution of QPOs during HIMS, was also independently ﬁtted with the POS
model and instantaneous shock location and compression ratios are calculated, since according to the POS
model, QPOs occur due to oscillation of the same shock(forms at the outer edge of the CENBOL) which we use
in TCAF model code. We also predicted frequency of theobserved type-C QPOs from TCAF model ﬁtted shock
parameters to compare those with observed and POS
model ﬁtted values, using the following relation,
3

νs0
c
1
νQP O = tinf
= 2GM
1/2 ] ,
BH [R Xs (Xs −1)
all

84

two (MAXI J1659-152 and MAXI J1836-194) are shorter
orbital period BHCs, where soft states were missing.
These short orbital period sources belongs to a special
category (type-II or harder-type) of sources (for e.g.,
XTE J1118+480, Swift J1753-0127, etc.) where accretion disks are dominated with the low angular momentum sub-Keplerian halo matters. During the outbursts,
Keplerian disk rate does not rises much compared to thehalo rate, such that it is able to cool down CENBOL
completely. On the other hand, MAXI J1543-564 is a
type-I or classical-type, where all spectral states are observed (see, Stiele et al. 2011).
Spectral ﬁts of these three MAXI objects with the
TCAF solution, provide us a direct evidence of two components accretion ﬂows as they are observed to vary independently during the outbursts. Spectral study with
the TCAF solution enables us not only to understandspectral properties of the BHCs but also to predict othertiming and physical parameters associated with the BHs.
Diﬀerent spectral states and their transitions from one
state to another are more clearly understandable fromthe variations of TCAF model ﬁtted physical ﬂow parameters and nature of QPOs. Frequency of the dominating
type-C QPOs also could be predicted from TCAF modelﬁtted shock parameters, since in TCAF we use same
shock, whose oscillation is believed to the fact of QPOs.
In ARRID, we saw that the correlation between spectral
and timing properties of transient BHCs is more fundamental than so-called ‘q’-diagram or HID. We also estimated probable mass range of the three MAXI sourcesfrom our TCAF model ﬁtted spectral analysis within agood accuracy. We also estimated X-ray jet ﬂuxes during
the 2011 outburst of MAXI J1836-194 from our spectral
study with the TCAF solution (see, proceeding of Janaet al. in this volume and Jana et al. 2017).
From our recent study of various BHCs with the
TCAF solution including three MAXI objects discussed
here, we can say that TCAF is the most accurate theoretical model to understand physics of accretion processes
around a black hole, which was our decade long conjecture. Spectral ﬁts with the model, not only explains
spectral properties and but also capable to explains various timing features including QPOs of BHCs. One canalso estimate physical source parameters, such as mass,
spin from spectral ﬁts with the TCAF model. Finally, it
is needed to mention that we require only six model input parameters including the mass of the black hole andnormalization to ﬁt a spectrum with the TCAF solution.
Acknowledgments
D.D. and D.C. acknowledge support from DST
sponsored Fast-track Young Scientist project fund
(SR/FTP/PS-188/2012). A.J. and D.D. acknowledges
support from ISRO sponsored RESPOND project fund

(ISRO/RES/2/388/2014-15).
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85

86

87

Inﬂow-Outﬂow Properties of Accretion Disk around MAXI J1836-194
with TCAF Solution during its 2011 Outburst
Arghajit Jana,1 Dipak Debnath,1 S. K. Chakrabarti,1,2 D. Chatterjee,1 A. A. Molla,1 , S. Mondal,3,1
1

Indian Centre for Space Physics, 43 Chalantika, Garia Station Road, Kolkata, 700084, India
2
SN Bose National Centre for Basic Science, Salt Lake, Kolkata, 700098, India
3
Instituto de Fı́sica y Astronomı́a, Universidad de Valparaı́so, Gran Bretana N 1111, Playa Ancha, Valparaso, Chile
E-mail: argha@csp.res.in; dipakcsp@gmail.com; chakraba@bose.res.in
Abstract
Galactic black hole candidate (BHC) MAXI J1836-194 was discovered by MAXI/GSC on August 29,
2011 at sky location RA = 18h 35m 43.43s , DEC = -19o 19’12.1”. We study the BHC during its 2011
outburst using archival data of RXTE/PCA with the TCAF Solution. From spectral ﬁts with the TCAF
model ﬁts ﬁle in XSPEC, we extracted physical accretion ﬂow parameters, such as, Keplerian disk rate,
sub-Keplerian halo rate, shock location and strength. Even probable mass of the black was also estimatedfrom our study. Quasi-periodic oscillations (QPOs) are observed sporadically during the outburst. Basedon the variation of ﬂow parameters and nature of QPOs, the entire outburst was classiﬁed into two spectral
states, hard and hard-intermediate. A deviation of the constancy of the TCAF model normalization wasobserved in ﬁve observations from hard-intermediate spectral states, where high radio-jets were present.
We also estimated jet X-ray ﬂux from our spectral study. Since, softer states were absent during theoutburst, we termed this outburst of MAXI J1836-194 as ‘failed-outburst’.
Key words: X-Rays:binaries – stars individual: (MAXI J1836-194) – stars:black holes – accretion,accretion disks – ISM: jets and outﬂows – radiation:dynamics

1. Introduction
Black hole candidate (BHC) MAXI J1836-194 was discovered simutaneously by MAXI/GSC and Swift/BAT
on August 29, 2011 at sky location RA = 18h 35m 43s .43,
DEC = -19◦ 19’12”.1. The BHC has short orbital period(∼ 4.7 hrs) and high spin parameters (a = 0.88 ± 0.03)(Jana et al., 2016 and references therin). The companion
is reported to be a high massive B[e] star (Cenko et al.
2011). After inclusion of two components advective ﬂow(TCAF) model in HeaSARC’s spectral analysis software
package XSPEC as a local additive table model, accretion ﬂow dynamics of several BHCs (for e.g., GX 339-4,
H 1743-322, MAXI J1659-152, MAXI J1543-564, etc.)
were studied and explained well (Debnath et al. 2014,2015a,b; Mondal et al. 2014, 2015, 2016; Jana et al.
2016; Chatterjee et al. 2016; Molla et al. 2016, 2017;
Bhattacharjee et al. 2017). We also estimated X-ray jet
outﬂow rates and mass of unknown BHCs from spectral
analysis using the TCAF solution.
2. Observation and Data Analysis
We use 2.5-25 keV RXTE/PCA archival data for 35 observation spread over the enire outburst. We carry out

our data analysis with TCAF Solution in XSPEC to extract physical ﬂow parameters, such as, Keplerian disk
rate (ṁd ), sub-Keplerian halo rate (ṁd ), shock location(Xs ), compression ratio (R) and mass of the BH. Low frequency quasi-periodic oscillations (QPOs) are observed
sporadically in power-density spectra during the entire
outburst.
TCAF normalization (N ) is a function of mass, distance and inclination angle of the disk. Thus constantvalue of N should be required to ﬁt all the data. However, higher values of N are required when jets arepresent, since in the present TCAF model, Jet calculations were not included. This gives us opportunity to
calculate jet X-ray ﬂux (Fjet ) from the total X-ray ﬂux(FX ). Then we freeze value of N at its lowest value (assuming there was no jet on that day) to calculate thedisk X-ray ﬂux (Fdsk ). Then we take diﬀerence of Fdsk
from FX to calculate Fjet as, Fjet = FX − Fdsk .
3. Result and Discussions
3.1. Timing Properties
Low frequency QPOs are observed sporadically on andoﬀ during entire 2011 outburst of MAXI J1836-194. Al-

88though a general trends of monotonically increasing frequencies in the rising phase and monotonically decreasing nature in the declining phase of the outburst are
observed. Highest value of QPO frequency 5.17 Hz was
observed on MJD=55820, when a transition between two(rising and declining) hard-intermediate spectral states
are observed. Non-satisfaction of the resonance condition between cooling and infall time scales (Chakrabarti
et al. 2015 and references therein) may be the reason
behind the non-observation of QPOs on daily basis.
3.2. Spectral Properties
Based on the nature of low frequency QPOs, accretion rate ratio (ARR = ṁh /ṁd ) and photon indices(Γ) we classiﬁed entire outburst into two harder spectral states: hard (HS) and hard-intermediate (HIMS),
which are observed in the sequence of: HS (Ris) → HIMS
(Ris) → HIMS (Dec) → HS (Dec). In the HS (Ris),
both ṁh and ṁd increases. ARR reaches on maximum
value of MJD=55808 when it enters to HIMS (Ris). In
this state, ṁd increases and reaches maximum value on
MJD=55820. ARR was minimum on this day. Highest
value of QPO was also observed in the same day. We
deﬁne this day as transition day between two HIMS. After that the BHC enters on HIMS (Dec). ṁd is observed
to be decreased but ṁh is observed to be increased. On
MJD = 55830, ARR was found to have local maxima as
the BHC enters in HS (Dec). The shock becomes strong
and moves away from the black hole during this state.
Both ṁh and ṁd are found to be decreased but ARR
increased. No sign of soft spectral state was found. Thismay be due to the BH is immersed in the excreation diskof the companion which is reported to be a B[e] star.
From the spectral study with TCAF solution, we predicted most probable range of the mass of the BHC as∼ 7.5 − 11 M . We also ﬁnd ṁh attains value maximum∼ 10 days before ṁd in both rising and declining phase.
This allow us to estimate the viscous time scale for thisoutburst is ∼ 10 days (see, Jana et al. 2016 for moredetails).
3.3. Estimation of X-ray Jet Flux
We ﬁnd TCAF normalization (N ) is varied over a small
range of ∼ 0.25 − 0.35 except ﬁve observations in hardintermediate state (around ∼ MJD = 55820), whenhigher values are required. But, according to the TCAF
solution, for a particular source, it should be constant,since N depends on intrinsic physical parameters (suchas, mass, distance, disk inclination angle) of the source.
The deviation could be due to non-inclusion of jet outﬂow processes in the current TCAF model ﬁts ﬁle. This
allows us to extract jet X-ray ﬂux contributions from total X-rays as deﬁned in §2. From the diﬀerences of FX
and Fdsk , we calculated the jet X-ray ﬂux (Fjet ). We
require highest value of normalization (N ∼ 1.68) on

MJD = 55820. We ﬁnd jet X-ray ﬂux (Fjet ) starts toincrease from MJD = 55810 and was maximum on MJD= 55820. After that it decreases. The variation of N
has similar trends as of X-ray jet ﬂuxes. Observed radio ﬂuxes (FR ) also showed roughly similar variations.
We also see about ∼ 25% X-ray is contributed by jet on
average. When jet is strongest, the contribution is upto
∼ 44%.
We also ﬁnd the correlation between FR and FX acb. Here, we obtain corcording to the relation FR ∝ FX
relation index as b = 1.85 ± 0.17. We also ﬁnd anothercorrelation between FR and Fjet . In this case, the correlation index is found to be as b = 0.97 ± 0.11. Our result
is consistent with Russel et al. (2015).
Acknoledgement
A.J. and D.D. acknowledge support from ISRO sponsored RESPOND project fund (ISRO/RES/2/388/201415). D.D. and D.C. also acknowledge support from
DST sponsored Fast-track Young Scientist project fund
(SR/FTP/PS-188/2012).
References
Cenko, S. B., Miller, A. A., & Bloom, J. S., et al., 2011,
ATel, 3614, 1
Bhattacharjee, A., et al., 2017, MNRAS, 466, 1372
Chakrabarti, S. K., & Titarchuk, L. G., 1995, ApJ, 455,
623
Chakrabarti, S. K., Mondal, S., & Debnath, D., 2015,
MNRAS, 452, 3451
Chatterjee, D., Debnath, D., & Chakrabarti, S. K., etal., 2016, ApJ, 827, 88
Debnath, D., Chakrabarti, S. K., Mondal, S., 2014,
MNRAS, 440, L121
Debnath, D., Mondal, S., & Chakrabarti, S. K., 2015a,
MNRAS, 447, 1984
Debnath, D., Molla, A. A., Chakrabarti, S. K., &Mondal, S., 2015b, ApJ, 803, 59
Jana, A., Debnath, D., & Chakrabarti, S. K., et al., 2016,
ApJ, 819, 107Jana, A., Chakrabarti, S. K., & Debnath, D., et al., 2017,
ApJ (submitted)
Molla, A. A., Debnath, D., & Chakrabarti, S. K., et al.,2016, MNRAS, 460, 3163
Molla, A. A., Chakrabarti, S. K., Debnath, D., &Mondal, S., 2017, ApJ, 834, 88
Mondal, S., Debnath, D., & Chakrabarti, S. K., 2014,
ApJ, 786, 4
Mondal, S., Chakrabarti, S.K., & Debnath, D., 2015,
ApJ, 798, 57
Mondal, S., Chakrabarti, S.K., & Debnath, D., 2016,
Ap&SS, 361, 309
Russel, T. D., Miller-Jones, J. C. A., & Curran, P. A.,
2015, MNRAS, 450,1745

89

Enhanced power spectrum analysis of short-term X-ray variability using
MAXI/GSC data
Tomofumi Kawase and Hitoshi Negoro

Department of Physics, Nihon University, 1-8 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8308, Japan
E-mail(TK): kawase@cygnus.phys.cst.nihon-u.ac.jp
Abstract
Power spectral densities (PSDs) of bright black hole candidates (BHCs) and neutron star X-ray binaries
must provide the information about, for instance, state transitions and quasi-periodic oscillations. However,a triangular window function of the eﬀective area of MAXI/GSC aﬀects the light curves of individual
sources due to the scanning observation. Suzuki (2015) established the method of evaluating resultant
PSDs correctly, and showed its usefulness for some BHCs. We modiﬁed his programs to analyze data more
eﬃciently. Scan duration to analyze in the previous method was ﬁxed to the minimum scan duration, 40 s.
We extended the duration to use data as long as possible (≤ 200 s). As a result, we could more correctly
evaluate a PSD with a break at a lower frequency. In addition, we also improved signal-to-noise ratios of
PSDs for not-so-bright sources by cutting the edges of the window function, and evaluated its eﬀects on
the PSDs analytically.
Key words: accretion, accretion disks — X-rays: binaries — stars: individual: Cyg X-1

1. Introduction
If very bright new transients appear, MAXI/GSC data
must be important. MAXI scans the whole sky every
92 minutes. Therefore, the probability that MAXI/GSC
can discover bright new transients is higher than otherobservatories. Moreover, although observatories focusing
X-rays can not observe extremely bright sources due todetector saturation, MAXI can still do that because it
does not have mirrors to focus X-rays.
For these reasons, we must prepare the method to analyze short-term variability using MAXI/GSC data forupcoming bright new transients.
2. Power spectral density, PSD, for MAXITime duration of a scan transit for a point source is 40200 s. It depends on the direction of the source orthogonal to the scan direction. In the previous method to
obtain a power spectral density (PSD) by Suzuki (2015),
the time duration to analyze was ﬁxed to the minimum
time duration, 40 s. We modiﬁed his programs to use
longer time duration data, which enabled us to deal withlower frequencies of the PSD. It is known that the PSD of
the black hole candidate (BHC) in the hard state shows
breaks at frequencies of 0.01-1 Hz (e.g., Pottschmidt etal. 2003). On the other hand, there is no break in thisfrequency band when the BHC is in the soft state (e.g.,

Cui et al. 1997). By this improvement, it becomes easier to ﬁnd the low frequency breaks, and to notice statetransitions not by energy spectra but by PSDs.
We calculate PSDs with reference to van del Klis
(1989), and normalize them according to Leahy et al.
(1983) and Miyamoto et al. (1994).
3. Eﬀects of a window function on the PSD
3.1. Time variability of the eﬀective area
The eﬀective area of GSC for a point source changes
sequentially by a collimator to restrict the direction ofincoming X-ray photons (ﬁgure 1). Therefore, a triangular window function aﬀects the light curve and its PSD.
To analyze the PSD using GSC data, we need to evaluate the eﬀect of the window function correctly.

Fig. 1. The movement of GSC in scanning observation and the time
variability of the eﬀective area for a point source. When the source
is in front of GSC, the eﬀective area reaches the maximum value.

1
0.1

Normalized Power4 10í3 0.01

í2

0

0

χ

χ

2

2

410í3 0.01

0.1

1

Normalized Power

10

10

100

100

90

0.01

−∞

Eq (2) is the Fourier transform of the window functions
in ﬁgure 2. T is the scan duration for a point source,and α is the cutting ratio for a new window function(0 < α ≤ 1). If α = 1, this equation corresponds to thewindow function of the left panel of ﬁgure 2.

iωT
2
αωT
2
W (ω) =
4 sin2eαω 2 T 2
4αωT+ωT (1 − α) sin
2


(2)

4. PSD of random shot-noise data
4.1. Break at a lower frequency
We performed simulations using random shot-noise datato conﬁrm whether it is possible to evaluate a break of
PSD at a lower frequency. A time constant of the shots
was assumed to be, for instance, 4 s. The lowest frequency of the PSD obtained by 40 s data was at 0.025
Hz. On the other hand, if we set the time duration for
100 s using our new programs, the lowest frequency was
at 0.01 Hz. We produced 48 light curves, which were atypical number of data sets for 3 days data of GSC. The
count rate of them corresponded to 400 mCrab.
Fitting the PSD by a Lorentzian model gives a breakfrequency at 0.0210 ± 0.0343 Hz. This does not tightly
constrain a break frequency because of the large error.
On the other hand, a break frequency of the PSD forthe longer time duration data was obtained at 0.0370 ±0.0029 Hz. This frequency is in good agreement withthe analytical break frequency of 0.0398 Hz for the time
constant of 4 s.

1Frequency [Hz]

10

0.01

0.1

1
Frequency [Hz]

10

0.1

1
Frequency [Hz]

100
10

0.1
0.01
10í5

10í4

10í3

Normalized Power

1

1
0.1
0.01
10í3

Normalized Power

Cyg_Xí1_Soft_State_MJD55500í55514

10

10

100

Cyg_Xí1_Hard_State_MJD55228í55243

10í4

3.2. New window function for not-so-bright sources
In case of not-so-bright sources, the background counts
unaﬀected by the triangular window function are relatively higher than the source signal at the beginning andthe end of the scan. Therefore, we adopt a new windowfunction that cuts the edges of the triangle to remove
low signal-to-noise ratio data (ﬁgure 2). The cutting ratio should be adjusted by the brightness of the source.
Eq (1) shows the relation of a source signal x(t), a
window function w(t), and their Fourier transforms X(ω)
and W (ω). The Fourier transform of the product of x(t)and w(t) is the convolution of X(ω) and W (ω), X ∗ W .
 ∞x(t)w(t)eiωt dt = X(ω) ∗ W (ω)
(1)

0.1

Fig. 3. A PSD of random shot-noise aﬀected by a triangular windowfunction (Left), and a new window function (Right, α = 0.5 inEq.2). Red plots are data, and blue plots are the eﬀect of thewindow function. The PSDs are ﬁtted by a Lorentzian.

10í5Fig. 2. Left: The previous triangular window function. Right: A new
window function for not-so-bright sources.

0.1

1
Frequency [Hz]10

Fig. 4. PSDs of Cyg X-1 in the hard state (Left, MJD 55228-55243,
2-20 keV), and the soft state (Right, MJD 55500-55514, 2-20
keV).

4.2. PSD using a new window function
Next, we conﬁrm that the PSD is correctly evaluated bya new window function. We prepared 100 light curvesof random shot-noise data for 100 s. The count rateof them corresponded to 400 mCrab. A time constantof the shots was 0.2 s. We also prepared light curves
removing both edges of them, and made a PSD (ﬁgure3). Since α was 0.5, the time duration of them was 50 s.
We could conﬁrm that a large eﬀect of the windowfunction appeared in the lowest frequency bin on both
PSDs, and evaluate both of them correctly.
5. PSD of Cyg X-1
We made PSDs of the BHC, Cyg X-1, for 15 days (ﬁgure4). We used only one camera data. A PSD of the hardstate was ﬁtted by the sum of two Lorentzians model,
and a PSD of the soft state was ﬁtted by a power-law
model. We conﬁrmed these results were consistent withprevious results (e.g., Pottschmidt et al. 2003).
6. Summary and Future
We modiﬁed the programs to obtain a PSD of short-term
variability using MAXI/GSC data, and conﬁrmed thesewere useful for the actual data. We are going to analyze not-so-bright sources using this new window function and its α optimized for the background counts.
References
Cui, W. et al. 1997, The Transparent Universe, 382, 209
Leahy, D.A. et al. 1983, ApJ 272, 256-258
Miyamoto, S. et al. 1994, ApJ 435, 398-406
Pottschmidt, K. et al. 2003, A&A, 407, 1039
Suzuki, K. 2015, Master’s thesis, Nihon University
van del Klis, M. 1989, in Timing Neutron star, NATOASI Series C, Vol. 262, p27-70

91

Reinvestigation of the relation between the state transition and the massaccretion rate in black hole candidates
Takahiro Masumitsu1 and Hitoshi Negoro1
1

Nihon University, 1-8-14 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8303, Japan
E-mail(TM): masumitsu@cygnus.phys.cst.nihon-u.ac.jp
Abstract

State transitions observed in black hole candidates (BHCs) are thought to result from changes in themass accretion rate. To estimate the distance to 6 BHCs discovered by MAXI, we verify the luminosityand mass-accretion-rate relation that the X-ray luminosity at the high/soft-state (HSS) to low/hard-state(LHS) transition is 1-4% of the Eddington luminosity, LEdd (e.g., Maccarone 2003). We also reinvestigate
two empirical relations in outburst, pointed out by Yu et al. (2007), that the LHS peak ﬂux at the
beginning of an outburst is nearly proportional to the time since the ﬂux peak in the previous outburst,
and that the following HSS peak ﬂux is nearly proportional to the LHS peak ﬂux. Using MAXI/GSC
and Swift/BAT data of GX 339-4, we conﬁrmed the constancy (0.01-0.03LEdd ) of the X-ray luminosity
at the HSS-to-LHS transition. The other two relations, however, could not be conﬁrmed in MAXI/GSC,
RXTE/ASM, and Swift/BAT data of GX 339-4 and H 1743-322.
Key words: accretion, accretion disks — stars: black hole candidates: individual (GX 339-4, H
1743-322)

2.

Observations and Data Reduction
To reexamine these relations, we analyzed GX 339-4
and H 1743-322. The distance and the mass of GX 339-4
we used are 6 kpc and 5.8 M  , respectively (Hynes etal. 2004).
2.1. Reevaluation of light curves
In order to determine more reliable time at the statetransition, we reevaluated MAXI light curves. Even if
the scanning observation terminated, for instance, toprotect the detectors, the lack of the background region

is not taken into account in the public data, resulting infake brightening. On the other hand, in the on-demanddata, such data are excluded in the default conﬁguration. To solve these problems, we used the image ﬁtmethod that took into account the point spread function of the detectors (Morii et al. 2016). As a result, we
obtained more continuous and more reliable light curves,especially for data of H 1743-322 near the Galactic center, inﬂuenced by X-ray emission from the Galactic plane(Fig. 1).
H 1743í322 (4í10 keV)counts/s/cm2

0.15

RIKEN Data Products

0.1
0.05
0
0.15

photons/s/cm2

Introduction
MAXI discovered 6 black hole candidates (BHCs), but
the distance to these new sources has not been known
yet. To estimate the distance, we reexamined the relation that the luminosity at the state transition fromthe high/soft state (HSS) to the low/hard state (LHS) is
1-4% of the Eddington luminosity (Maccarone 2003) or
0.1-10% of the Eddington luminosity (Dunn et al. 2010).
We also investigated the following two empirical relations with the mass accretion; the nearly proportional
relation between the LHS peak ﬂux and the time since
the LHS peak in an previous outburst, and the proportional relation between the LHS and HSS peak ﬂuxes in
the outburst (Yu et al. 2007).

RIKEN Ondemand

0.1
0.05
0
0.15

photons/s/cm2

1.

Image Fitting Method

0.1
0.05
0
5.69×104

5.7×104

5.71×104
Time (MJD)

5.72×104

Fig. 1: Light curves of H 1743-322 (2015/02/05-08/04)(from the top, Riken public data, Riken on-demand data,and image ﬁtting method data)

92

counts/s/cm2

MAXI/GSC(4.0í10.0 keV)

0.4
0.2
0
0.15

counts/s/cm2

Swift/BAT(15.0í50.0 keV)

0.1

4í10 keV Xíray Flux (Crab)

GX339í4: Light Curve
0.6

0.2

1.5

Determination of the time at the state transition
4í10 keV Xíray Flux (Crab)

2.2.

1

0.5

0

0

500

1000

1500

0.15

0.1

0.05

0

0

200

Time (Day)400

Time (Day)0.05

Fig. 4: Peak ﬂuxes in individual outbursts vs. the timeintervals from the previous outburst in GX 339-4 (leftpanel) and H 1743-322 (right panel) (red: MAXI/GSC,
black: RXTE/ASM)

0
1
Ratio(15.0í50.0 keV/4.0í10.0 keV)

0.5
05.52×104

5.54×104

5.56×104

5.58×104

MJD

Fig. 2: GX 339-4 (09/08/15-11/08/27) (from the top,
MAXI/GSC data 4-10 keV, Swift/BAT data 15-50 keV,and the hardness ratios (BAT/GSC))

3.3. Is the LHS peak ﬂux of the outburst nearly proportional to the HSS peak ﬂux?

By combining MAXI/GSC and Swift/BAT data
(http://swift.gsfc.nasa.gov/results/transients/weak/GX339-4/), we can estimate the time at the state transition
more accurately (Fig. 2). We regard the peak-ﬂux timein the Swift/BAT data around the transition, seen in
the hardness ratios, as the LHS peak-ﬂux time at thestate transition.
3. Analysis
3.1. Is HSS-to-LHS transition luminosity 1-4% of the Eddington luminosity?

2í200 keV Peak HSS Luminosity(1037ergs/s)

20

15

10

5

0

0

5

10

15

20

2í200 keV Peak LHS Luminosity(1037ergs/s)

Fig. 5: Relation between LHS peak luminosities and HSS
peak luminosities in GX 339-4 (black: in 2011, red: in2015. The two diﬀerent luminosities in 2015 are obtainedby two acceptable spectral models.)
We used the MAXI/GSC data of GX 339-4. The LH
peak ﬂuxes (luminosities) in the two outbursts were different by 1.5-2.0 times, but the HS peak ﬂuxes are nearlysame (1.05) (Fig. 5).

(a) 2011/02/06-10(b) 2015/08/16-20

Fig. 3: Energy spectra during HSS-to-LHS transitions in
GX 339-4
We examined two luminosities at the HSS-to-LHS
transitions using MAXI/GSC spectra of GX 339-4, andconﬁrmed that both values were in the range of 1-3% of
the Eddington luminosity (Fig. 3).
Is the peak ﬂux of the outburst nearly proportional to
the time since the ﬂux peak in the previous outburst?
We analyzed MAXI/GSC and RXTE/ASM data of
GX 339-4 over the past 20 years, and MAXI/GSC and
Swift/BAT data of H 1743-322 over the past 7 years. We
could not conﬁrmed the proportional relation between
the LHS peak ﬂux and the time since the peak in theprevious outburst (Fig. 4).

4. Summay/Tasks
We conﬁrmed the constancy of the X-ray luminosity at
the HSS-to-LHS transition in GX 339-4 as was expectedfrom a simple accretion disk theory. On the other hand,
the two kinds of the linear relations between the LHS
and HSS peak-ﬂuxes and the time interval were not conﬁrmed not only in GX 339-4 but in H 1743-322. Finally,
we note that out of 9 outbursts in H 1743-322 over the
past 7 years, 4 outbursts did not show the state transition. We are investigating the cause of the diﬀerence.
Detailed analysis results here will be reported elsewhere.

3.2.

References
Maccarone 2003, A&A, 409, 697
Yu et al. 2007, ApJ, 663, 1309
Dunn et al. 2010, MNRAS, 403, 61
Hynes et al. 2004, ApJ, 609, 317
Morii et al. 2016, PASJ, 68, SP1, a11Nakahira et al. 2013, J. Space Sci. Info. Jpn., 2. 29

93

A model for the ‘outlier’ radio/X-ray correlation in black hole X-ray
binaries
Erlin Qiao 1,2
Key Laboratory of Space Astronomy and Technology,
National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China
2
School of Astronomy and Space Sciences, University of Chinese Academy of Sciences,
19A Yuquan Road, Beijing 100049, China
E-mail:qiaoel@nao.cas.cn
1

Abstract
Observationally, it is found that there is a strong correlation between the radio luminosity and the X-ray luminosity in black hole X-ray binaries. The ‘universal’ correlation, i.e., LR ∝ L∼0.6
was well explained by the modelXof the coupling of a radiatively ineﬃcient accretion ﬂow and a jet. Recently a growing number of sources havebeen found with a complicated radio/X-ray correlation, e.g., LR ∝ L∼1.4
for LX /LEdd  10−3 , which is often calledX‘outlier’ correlation. In this paper, we interpreted such a ‘outlier’ radio/X-ray correlation within the frameworkof a disk corona-jet model, in which the matter in the accretion ﬂow and the matter in the jet are connected by aparameter, η, describing the fraction of the matter in the accretion ﬂow ejected outward to form the jet. By ﬁtting
the spectrum, it is found that the X-ray emission is dominated by the corona and the radio emission is dominated by
the jet. Meanwhile, we found that the value of η for the ‘outlier’ correlation for LX /LEdd > 10−3 is systematicallyless than that of the ‘standard’ correlation, which is consistent with the general idea that the jet is often relatively
suppressed at the high luminosity phase in black hole X-ray binaries.
Key words: accretion discs — black hole physics — X-rays — radio emission — H1743-322

1.

Introduction

Recently, a growing number of black hole X-ray binaries(BHXBs) have been discovered with a ‘outlier’ radio/X-ray
correlation of LR ∝ L∼1.4
for LX  10−3 LEdd , which is sugX
gested to be explained within the framework of the couplingof a radiatively eﬃcient accretion ﬂow and a jet (Corbel et al.
2004; Rodriguez et al. 2007; Soleri et al. 2010; Coriat et al.
2011; Corbel et al. 2013). Observationally, there is evidence
for the coupling of a disk corona and a jet at a high mass accretion rate and the coupling of a radiatively ineﬃcient accretion ﬂow (RIAF) and a jet at a low mass accretion rates (Wu et
al. 2013). Meanwhile, theoretically, for Ṁ  α2 ṀEdd (with α
the viscosity parameter, ṀEdd = 1.39 × 1018 M/M g s−1 ), theaccretion ﬂow will transit from a RIAF to a disk corona system (Qiao & Liu 2009, 2010, 2013; Narayan & Yi 1995b).
In this work, we proposed a disk corona-jet model to explain the ‘outlier’ radio/X-ray correlation of LR ∝ L∼1.4
forX
LX /LEdd  10−3 . We brieﬂy introduce the disk corona-jet
model and the results in Section 2. Section 3 is the conclusion.
2.

The model

In the model, the disk and corona are radiatively and dynamically coupled (Liu et al. 2002,2003; Qiao & Liu 2015). Thus,

the energy fraction dissipated respectively in the corona anddisk, and coronal density and temperature can all be selfconsistently determined for given black hole mass and accretion rate. Then the spectrum emitted by the disk and corona
can be calculated by Monte Carlo simulation. Assuming afraction of matter in the accretion ﬂow, η ≡ Ṁ jet / Ṁ, is ejected
outward to form the jet, we can also calculate the emergent
spectrum from the jet. So far, the theoretical understandingof the jet formation is poor. Speciﬁcally, it is diﬃcult to put
constraints on the dependence of η on Ṁ in our model, so weset η as an independent parameter on Ṁ to ﬁt the observations.

We calculate LR and LX at diﬀerent Ṁ, adjusting η to ﬁt theobserved ‘outlier’ radio/X-ray correlation of the black hole Xray transient H1743-322 for LX /LEdd > 10−3 . It is found thatalways the X-ray emission is dominated by the corona and theradio emission is dominated by the jet. See Figure 1 for thedetails. We noted that the value of η for the radio/X-ray correlation of LR ∝ L∼1.4
for LX /LEdd > 10−3 , is systematicallyX
less than that of the case for LX /LEdd < 10−3 , which is consistent with the general idea of jets often relatively suppressed at
the high luminosity phase in black hole X-ray binaries.

94

Fig. 1.

Left-hand panel: Emergent spectra of the disc corona-jet model around a black hole with M = 10M assuming α = 0.3 formodeling the radio/X-ray correlation of the black hole X-ray transient H1743-322 for L3−9keV / LEdd > 10−3 . From the bottom
up, the solid lines are the combined emergent spectra of the disc corona-jet model for Ṁ = 0.02, 0.05, 0.1, 0.3 and 0.5 ṀEdd ,and the corresponding dotted lines are the emergent spectra from the jet with η = 0.54%, 0.52%, 0.57%, 0.62% and 0.62%
respectively. Right-hand panel: L8.5GHz /LEdd as a function of L3−9keV /LEdd . The red ‘•’ are the observations for H1743-322,
the green ’are the observations for GX 339-4, and orange ‘’ are the observations for V404 Cyg. The dotted line is thebest-ﬁtting linear regression of H1743-322 for L3−9keV / LEdd > 10−3 . The dashed line is the best-ﬁtting linear regression of
GX 339-4 and V404 Cyg. The thick solid line is the model line, and the model spectra are shown in the left panel.
3.

Conclusions

We investigate the ‘outlier’ radio/X-ray correlation of LR ∝L∼1.4
for LX /LEdd  10−3 within the framework of a diskXcorona-jet model. We note an interesting result, i.e., the fraction of the ejected matter η (∼ 0.57%) for the radio/X-ray
correlation of LR ∝ L∼1.4
for LX /LEdd  10−3 , is systematiX
cally less than that of the case for LX /LEdd < 10−3 (at leastη  1%), which may put some constraints on the jet formation, i.e., by suggesting that the strength of the jet power
is relatively suppressed during the high luminosity phase in
BHXBs.
Recently, it is found that there exits a ‘outlier’ radio/X-ray
correlation in luminous active galactic nuclei (AGNs) (Donget al. 2014; Panessa et al. 2015). We suggested that such asteep radio/X-ray correlation in AGNs can also be explainedwith the disk corona-jet model, which may imply the similarities of the theory of black hole accretion across the scalesof several orders of magnitude. The study of the ‘outlier’radio/X-ray correlation in AGNs will be conducted in the future.

References

Coriat, M. et al., 2011, MNRAS, 414, 677
Corbel S., Coriat M., Brocksopp C., Tzioumis A. K., Fender
R. P., Tomsick J. A., Buxton M. M., Bailyn C. D., 2013,

MNRAS, 428, 2500
Corbel S., Fender R. P., Tomsick J. A., Tzioumis A. K.,
Tingay S., 2004, ApJ, 617, 1272
Dong A.-J., Wu Q., Cao X.-F., 2014, ApJ, 787, L20
Liu, B. F., Mineshige, S., & Shibata K., 2002, ApJ, 572, L17
Liu, B. F., Mineshige, S., & Ohsuga K., 2003, ApJ, 587, 571
Narayan, R., & Yi, I., 1995b, ApJ, 452, 710
Panessa, F. et al., 2015, MNRAS, 447, 1289
Qiao, Erlin, & Liu, B. F., 2009, PASJ, 61, 403
Qiao, Erlin, & Liu, B. F., 2010, PASJ, 62, 661
Qiao, Erlin, & Liu, B. F., 2013, ApJ, 764, 2
Qiao, Erlin, & Liu, B. F., 2015, MNRAS, 448, 1099
Rodriguez J., Cadolle Bel M., Tomsick J. A., Corbel S.,
Brocksopp C., Paizis A., Shaw S. E., Bodaghee A., 2007,
ApJ, 655, L97
Soleri P. et al., 2010, MNRAS, 406, 1471
Wu, Q., Cao, X., Ho, Luis C., & Wang, Ding-Xiong, 2013,
ApJ, 770, 31

95Suzaku Observation of the Eastern X-ray Lobe of SS 433
Hideki Uchiyama,1 Yuta Kanbe2
1

Faculty of Education, Shizuoka University, 836 Ohya, Suruga-ku, Shizuoka, 422-8529, Japan
Faculty of Science, Shizuoka University, 836 Ohya, Suruga-ku, Shizuoka, 422-8529, Japan
E-mail(HU): uchiyama.hideki@shizuoka.ac.jp

2

Abstract
We observed the eastern X-ray robe of SS 433 with Suzaku. The non-thermal component of the Xray spectrum cannot be described by a simple power-law model, but requires a cutoﬀ around 8 keV. We
also found that the cutoﬀ energy decreases as the distance from SS 433 increases. It is probably due tosynchrotron cooling of electrons accelerated near SS 433. Based on this, the magnetic ﬁeld around thelobe B is limited to be weaker than ∼ 75 μG. The high-energy electrons traveled a distance of ∼9.6 pc
among 650 yr if B =10 μG. The bulk velocity of the jet, which is 0.26 c near the compact object, keeps∼ 0.05c after the jet traveled a distance of ∼ 35 pc.
Key words: ISM: jets and outﬂows — X-rays: individuals:(SS433) — X-rays: individuals:(W50)

5.0



4.9

 4.8

2. Spectrum from the Entire Field of View
We made a spectrum from the entire ﬁeld of view of
XIS. It is shown in ﬁgure 2. Models of the Galaxy
background (Uchiyama et al. 2013) and the cosmic Xray background (CXB, Kushino et al. 2002) were added
for ﬁtting. As reported by Brinkmann et al. (2007), thespectrum of the X-ray lobe consists of a optically-thinthermal plasma component (APEC, Smith et al. 2001)
and a non-thermal component with interstellar absorption.
If we used a simple power law as the non-thermal
component, large residual appeared above 3 keV. Whenthe non-thermal component is a cutoﬀ power law with






1. Introduction
SS 433 is an X-ray binary and its compact object is oneof the most famous blackhole candidates. Jets with avelocity of 0.26 c exist near the compact object (Margon 1984). A radio-shell supernova remnant (W50) is
associated with SS433. The shell is extended toward theeast and west. It is thought that the interaction with thejets caused the extended structure. Diﬀuse X-ray emission is also associated with the extended radio shell. On
the east side, there is a bright structure like a knot (e.g.
Yamauchi et al. 1994).
To reveal the detail structure of the X-ray emission, we
observed the eastern X-ray lobe for 107 ks with Suzaku
XIS whose features are large eﬀective area and a low
stable non-X-ray background. The observed position isabout 35 arcmin away from SS 433. The X-ray image isshown in ﬁgure 1.

    
288.7

288.6



  



288.5

288.4



Fig. 1. X-ray Image of the observed region with Suzaku XIS. The
regions from which spatially divided spectra are extracted are also
shown (see section 3.).

%*$)((!

96%$)'#"
%#&%$$)*)%.&%+'"+

'#"%#&%$$)",!'%*$

 #&"&%+'"+)%"''%'






/ %




 

 

*)%.&%+'"+/ %     



Fig. 2. Spectrum from the entire ﬁeld of view. The residuals of theﬁtting with simple or cutoﬀ power-law models for the non-thermal
components are also shown.

hνcutof f ∼ 8 keV, the residual disappeared and reduced
χ2 also became much smaller (ﬁgure 2). We, for the ﬁrst
time, found that the non-thermal component requires acutoﬀ.
3. Spatial Proﬁle of the Cutoﬀ Energy
We made spectra from the regions of A, B, C and D
shown in ﬁgure 1 to study the spatial proﬁle of the cutoﬀ energy hνcutof f . The regions of A, C and D are annuluses centered at SS 433. The region B corresponds
to the bright knot-like structure but its spectrum is notmuch diﬀerent from those of the other region. We ﬁtted
4 spectra simultaneously. The temperature of the thermal component, the photon index of the non-thermal
component and interstellar absorption were linked for all
the regions. The normalizations of the components andhνcutof f of determined independently for each region.
The obtained spatial proﬁle of hνcutof f is ﬁgure 3. Thehνcutof f decreases as the distance from SS 433 increases.
The spatial dependence of photon index was reported byprevious studies (Yamauchi et al. 1994, Brinkmann etal. 2007) but we found that it can be explained by thechange of hνcutof f .
4. Discussion
The probable origin of the non-thermal component is
synchrotron radiation of high-energy electrons. The cutoﬀ is probably due to the synchrotron cooling of the elec-

Fig. 3. Spatial proﬁle of hνcutof f obtained by simultaneous ﬁttingof spectra from region A–D.

trons. The higher the electron energy, the shorter thecooling time. The spatial change of the cutoﬀ hνcutof fshows that the electrons accelerated near SS 433 moveoutward while cooling. Our result suggests that theacceleration occurs outside of the ﬁeld of view of thisSuzaku observation.
Figure 3 indicates that the electrons traveled the distance between A and C, 9.6 pc (the distance between the
Sun and SS 433 is 5.5 kpc; Blundell & Bowler 2004), in
the cooling time of electron which radiate photons withthe energy of hνcutof f ∼ 5.5 keV typically. The upper
limit of the magnetic ﬁeld B near the eastern X-ray lobe
is obtained from the constraint that the speed of electrons is slower than the light velocity. Therefore, B is
weaker than 74μ G. A similar calculation between C andD gave a limit of B < 79μ G.
Hayashi et al. (2009) reported that the lower limit of
the interstellar magnetic ﬁeld in the western X-ray lobe
is 6.8 μG from the TeV γ-ray observation. Considering
this and the upper limit obtained above, we assumed that
B=10 G. If there is no reacceleration, the electrons arethought to be traveled 9.6 pc during ∼ 600 yr (A–C) or∼ 700 yr (C–D). This result shows that the bulk velocityof the jet keeps ∼ 0.05 c after the jet traveled a distanceof ∼ 35 pc. This is close to a expected velocity after
standing shock of the velocity of 0.26 c, 0.26/4 = 0.063 c.
References
Blundell, K. M., et al. 2004, ApJL, 616, L159
Brinkmann, W., et al. 2007, A&A, 463, 611
Hayashi, S., et al. 2009, Astroparticle Physics, 32, 112
Kushino, A., et al. 2002, PASJ, 54, 327
Margon, B. 1984, ARA&A, 22, 507
Smith, R. K., et al. 2001, ApJL, 556, L91
Uchiyama, et al. 2013, PASJ, 65, 19
Yamauchi, S., et al. 1994, PASJ, 46, L109

97

The broad-band X-ray power spectrum in black hole binaries: frommonths to milliseconds
Magnus Axelsson1
1

Tokyo Metropolitan University
E-mail(MA): magnus@tmu.ac.jp
Abstract

The X-ray emission from black hole binaries shows variability on timescales from milliseconds to years.
While the power spectrum has been well characterised above ∼0.01 Hz for a number of sources, thefrequency range below 10−5 Hz is much less studied. The monitoring capabilities of MAXI make it aunique tool in such studies. Combined with RXTE observations, we compare the low and high frequency
power spectrum in hard and soft states of bright sources such as GX 339−4 and Cygnus X-1. We ﬁnd thebroad-band variability to be more stable across states at low frequencies, and discuss our results in thecontext of propagating ﬂuctuation models.

98

X-ray monitering of the Galactic microquasar V4641 SgrKazutaka Yamaoka1
1

Nagoya University
E-mail(KY): yamaoka@isee.nagoya-u.ac.jp
Abstract
V4641 Sgr is the black hole (BH) candidate which exhibits a giant outburst with an X-ray intensity
up to 12 Crab on September 1999. Since then, several small outbursts with 1–100 mCrab level havebeen observed almost every year. However, it remains unclear whether this source has typical spectral
states such low/hard and high/soft states seen in BH candidates and what X-ray emission mechanisms are.
Large-scale radio jets have also been observed from this source, hence it is very important source to study
the relation between jets and accretion ﬂows. Since 2008, we have monitored this source in an outburstwith MAXI/GSC and Swift/XRT, and carried out detailed observations with Suzaku and NuSTAR in2014. The observed X-ray ﬂux ranges up to 40 mCrab. The spectra of this source look like high/soft state,
and can be well explained by multi-color disk model with an innermost temperature of 1.0–1.5 keV, but
the innermost radii, 8–21% of Schwarzchild radii, are too small to explain its emission region and does notkeep constant. We also detected iron emission lines at 6.8 keV due to highly ionized irons from Swift/XRT,NuSTAR, and Suzaku spectra, and ﬁrst revealed the presence of photo-ionized plasma in this source. In
this paper, we will discuss the origin of X-ray continuum and iron emission lines, and similarity to thepeculiar microquasar V404 Cyg.

99

100

101

Low
Mass
X-ray
Binaries
andSuperbursts

102

103

Spectral states in NS-LMXBs observed with MAXI/GSC and Swift/BAT
Kazumi Asai,1 Tatehiro Mihara,1 Masaru Matsuoka,1 and Mutsumi Sugizaki1
1

MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198
E-mail(KA): kazumi@crab.riken.jp
Abstract

We review the spectral states in NS-LMXBs observed with MAXI/GSC and Swift/BAT. The spectral
state transitions (soft/hard) were seen in the NS-LMXB transients. This was caused by the instability in astandard accretion disk. Sensitive observations with GSC and BAT showed outbursts even with small peakluminosities and shorter durations. Such “mini-outbursts” were observed in four sources; 4U 1636−536,
4U 1705−44, 4U 1608−52, and GS 1826−238 (Asai et al. 2015). We understand it as a “purr-type” ofdisk instability predicted by Mineshige and Osaki (1985), whereas the large normal outburst is named
as a roar-type. There are several atoll NS-LMXBs, which are persistent and always in the soft state.
Although the hardness ratio (HR) in GSC and BAT observations is always low (i.e. soft state), the HR
distributed in two (higher/lower) groups. The higher group contains 4U 1820−30 and 4U 1735−44, while
the lower group contains GX 3+1, GX 9+1, GX 13+1, and GX 9+9. We interpreted the diﬀerence tocome from the size of the Comptonized cloud. The high HR was resulted from a large Comptonized cloud,which might be dammed at the Alfvén radius by rather high magnetic ﬁelds (Asai et al. 2016). In thebrightest NS-LMXB (Z sources), there are three branches; Horizontal Branch (HB), Normal Branch (NB)
and Flaring Branch (FB). Although the diﬀerences of HB and NB in the spectral and timing properties
indicated some changes in the geometrical structure, the physical understanding is still controversial. We
suggest that the diﬀerence of HB and NB might be explained by existing of the disk evaporation.
Key words: accretion, accretion disks — X-rays:neutron — X-rays:binaries

1.

Introduction

Low mass X-ray binaries with a weakly magnetized neutron star (NS-LMXB) are divided into two main subclasses (Z sources and Atoll sources) based on their behavior on the color–color the diagram and hardness–
intensity diagram (Hasinger & van der Klis 1989). Z
sources are generally bright and sometimes become close
to the Eddington luminosity (LE ). On the other hand,
Atoll sources are generally less bright, L ≤ 0.5 LE . The
NS-LMXBs show two kinds of X-ray spectra depending on their X-ray luminosities. In the high luminosity(> 1037 erg/s ) it is dominated by blackbody components,
and called the soft state. On the other hand, in the low
luminosity( < 5×1036 erg/s ) it is dominated by a Comptonized component, and called the hard state. Z sources
are usually in soft state, whereas some of Atoll sources
exhibit spectral state transitions (soft or hard).
The present study focuses on the properties of thespectral state of Z sources and Atoll sources. The analyzed data are MAXI (Matsuoka et al. 2009)/GSC (Gas

Slit Camera: Mihara et al. 2011; Sugizaki et al. 2011) 1
and Swift (Gehrels et al. 2004)/BAT (Burst Alert Telescope: Barthelmy et al. 2005) 2 . The obtained countrates of GSC and BAT were converted to luminosities
by assuming a Crab-like spectrum (Kirsch et al. 2005)
and the distances listed in table 1. The assumption of
Crab-like spectrum is acceptable in the hard state, because the energy spectrum is dominated by the Comptonized emission approximated by a power law with thephoton index of 1–2. On the other hand, in the soft
state, the energy spectrum is dominated by the thermal
emission. The obtained luminosity by assuming Crablike spectrum is underestimated in the 2–10 keV band,
but is overestimated in the 15–50 keV band. As a result,the hardness ratio (HR) of two energy bands is overestimated by 2.0 times (see Asai et al. 2015 for detail).
Therefore, we handle only the relative diﬀerence.
In this paper, ﬁrst, we introduce two kinds of spectral state transitions, normal outburst and mini outburst*1 <http://maxi.riken.jp/>.
*2 <http://heasarc.gsfc.nasa.gov/docs/swift/results/transients/>.

(Asai et al. 2015). Next, we show two groups in a soft
state of Atoll sources, and then propose that the diﬀerence would come from the surface magnetic ﬁeld of aneutron star (Asai et al. 2016). Finally, we discuss thediﬀerence between Horizontal and Normal branch (HB
and NB) of Z sources. We suggest that the diﬀerence
might be explained by the disk evaporation.
Table 1. Distances used in this paper

Name
Atoll sources
Aql X-1
4U 1608−52
4U 1636−536
4U 1705−44
4U 1735−44
4U 1820−30
GX 3+1
GX 9+1
GX 9+9GX 13+1
GS 1826−238
XTE J1709−267
Z sources
Cyg X-2
GX 17+2
GX 340+0
Sco X-1
GX 5−1
GX 349+2

Distance (kpc)

Ref.∗

5
4.1
6
7.4
8.5
7.6
4.5
5
5
7
7
8.8

(1)
(2)
(1)
(1)
(2)
(1)
(3)
(1)
(4)
(1)
(1)
(1)

7.2
12.6
11
2.8
9.2
5

(1)
(5)
(4)
(1)
(4)
(4)

∗ (1) Liu et al. (2007), (2) Galloway et al. (2008), (3) Kuulkers and van der Klis (2000), (4) Christian and Swank(1997), (5) Lin et al. (2012).
2. Two kinds of spectral state transitions
We deﬁne the spectral state with MAXI/GSC (2–10 keV)and Swift/BAT (15–50 keV). First, we make the histograms (number of days) of the BAT/GSC hardness
ratios (HRs) of the seven NS-LMXBs (4U 1636−536,
4U 1705−44, 4U 1608−52, GS 1826−238, Aql X-1,
XTE J1709−267, and 4U 1820−30). We basicallyadopted the HR of the lowest point between the two
peaks as the threshold between the soft and hard states(see table 2 in Asai et al. 2015). Figure 1 shows the GSC
light curve, BAT light curve, and HRs (BAT/GSC) of
4U 1705−44. In this ﬁgure, the threshold separate thedata points into soft and hard states.
From the light curve, we noticed that the soft states
are separated into two luminosity classes. Comparing
with the average luminosity of soft state, one is above

BAT/GSC LBAT(erg sí1) LGSC(erg sí1)

104

4U 1705í44

1038
1037
1036
35
1037
10

1036
1035
1
0.1
0.01
55000

55200

55400

55600

55800

56000
MJD

56200

56400

56600

56800

Fig. 1. One-day light curves (LC) of 4U 1705−44 shown as an example. Top: GSC LC in the 2–10 keV band. Middle: BAT LC in the
15–50 keV band. Bottom: the hardness ratios (HRs) (BAT/GSC).
Vertical error bars represent 1 σ statistical uncertainty. Circles andtriangles represent data of soft and hard states, respectively. See
Asai et al. (2015) for other sources.

ten to 1037 erg s−1 , the other is below this value. So, we
calculated the average 2–10 keV luminosities of the softand hard states using the thresholds of the HR for sevensources. The soft states of all seven sources are separated
into two luminosity classes. If the average luminosity inthe soft state is below ∼ 1037 erg s−1 , the X-ray variability (repeated small increase) is called “mini-outbursts,”
whereas outburst with average luminosity above ∼ 1037
erg s−1 is called “normal outburst.” The mini-outbursts
were characterized by smaller amplitudes and shorter
duration than those of normal outbursts. Such minioutbursts were observed in the four sources (4U 1636–
536, 4U 1705–44, 4U 1608–52, and GS 1826–238). We
consider that the light curve of mini outburst is similar
to the calculated light curve for a “purr-type” of disk
instability predicted by Mineshige and Osaki (1985).
Normal outbursts in NS-LMXBs are usually interpreted by the disk instability model like the outburstof dwarf novae. In this model, local disk instabilities(S-shaped curves in the surface density and mass accretion rate diagram) are caused by the partial ionizationof hydrogen and helium. The theoretical disk instability model of Mineshige and Osaki (1983, 1985) predicts
both large-amplitude outbursts (roar type) and smallamplitude variability (purr type). The former is considered to correspond to normal outburst, but no correspondence has been found for the latter. Here, we proposethat purr-type outbursts correspond to mini-outbursts,
because the characteristics (small-amplitude and shortduration) of both types are quite similar.
The main diﬀerence between two types (roar type andpurr type) is α parameter of the standard disk. The αparameter is usually considered to depend on the disk
temperature. In order to explain a normal outburst,
two S-curves with a diﬀerent α values are used. As aresult, the width of the middle branch of S-curve is extended. This means that the instability can propagate
through the whole disk, and then outburst occurs. Thisis a roar type. In the case of one S-curve, the width

105

Luminosity (1036 erg/s)

100

(a)

(b)

(c)

4U182010

GS1826
4U1636

1
109

1010
1011
Disk radius (cm)Aql Xí1
4U16081012

Fig. 2. Average X-ray luminosity in the 2–10 keV band as a functionof outer disk radius. Open circles and ﬁlled triangles represent
data of soft and hard states, respectively. Theoretical lines of (a),(b), and (c) are calculated by equation (53) of Tuchman et al.
(1990) (C = 10−4 ). The irradiation temperatures of (a), (b),and (c) are Tirr = 10000 K, 5000 K, and 3000 K, respectively.

of the middle branch is narrow, and then the instability
propagates only a part of the disk. This is a purr type.
We considered that mini outburst occurs when the middle branch is narrow. Tuchman, Mineshige, and Wheeler(1990) reported that the middle branch can become narrow, when the disk is irradiated by the central sources.
So, we investigate the status of the outer disk, assumingthat the disk temperature is determined by X-ray irradiation. Figure 2 shows the average X-ray luminosities
and the outer radius of the accretion disk. Here, theouter radius was assumed as 0.35 of the binary separation (Smak 1982, and see Asai et al. 2015 in detail). For
4U 1820−30, we interpreted the spectral state transitionas other type because average luminosity in both soft andhard state are above ∼ 1037 erg s−1 (see Asai et al. 2015
in detail). The left side of (a) indicates a stable region,
where S-shape disappears. The right side of (c) is an
unstable region, where a normal outburst occurs. In theregion between (a) and (b), we observed mini-outbursts.
Therefore, we interpreted mini-outbursts were caused by
X-ray irradiation.
3. Hard tail emission in Atoll sources
Average hard-tail X-ray emission in the soft state of ninebright Atoll source (Aql X-1, 4U1608−52, 4U1705−44,
GX 3+1, GX 9+9, GX 13+1, GX 9+1, 4U 1735−44,
4U 1820−30) are investigated by using the light curves of
MAXI/GSC and Swift/BAT. Two sources (4U 1820−30
and 4U 1735−44) exhibit large HR (15–50 keV/2–
10 keV: HR > 0.1), while the other sources distributeat HR ≤ 0.1. In either case, HR does not depend on
the 2–10 keV luminosity. Therefore the diﬀerence of
HR is due to the 15–50 keV luminosity, which is Comp-

tonized emission. The Compton cloud is assumed to bearound the neutron star. In general, Comptonized emission is considered to be aﬀected by the inclination, originof seed photon, and electron temperature of Compton
cloud. However, we cannot see a clear diﬀerence of any
parameters between two groups from the results by several authors (see Asai et al. 2016). So, we propose thatthe diﬀerence between the two groups would be due tothe size of the Compton cloud. The group with high HR
would have large Compton cloud. Next, we consider thelocation of Compton cloud.
In the soft state, the location for Compton cloud iscontroversial. Here, let us take into account of the magnetic ﬁeld of the neutron star. The magnetic ﬁeld of aneutron star in LMXB is considered to be weak. However, in some cases, the eﬀect of the magnetic ﬁeld istaken into account. In the soft state, most of the gas in
the disk accretes onto the neutron star, because the disk
can extend to the neutron star or ISCO. The accretionﬂow is thermalized on the surface, and the emission fromthere is Comptonized by the plasma around the neutronstar. However, if the magnetic ﬁeld of the neutron staris relatively strong, the Alfvén radius would be larger
than the ISCO. In this case, the accretion ﬂow would be
stopped and spread around the Alfvén radius, and thenthe relatively large Compton cloud would be created.
We calculated the Alfvén radius as a function of luminosity (see ﬁgure 7 in Asai et al. 2016). Here, we adopted
a gas pressure dominant accretion disk. When L is1037 erg s−1 , RA is larger than ISCO for B > 1.5×108 G.
In this case, the accretion ﬂow would be stopped andspread at the Alfvén radius, and then the relatively large
Compton cloud would be created. As a result, the HR
would be large (for example HR > 0.09). On the other
hand, when B < 1 × 108 G, the RA is smaller than both
ISCO and NS surface. In this case, the accretion ﬂow is
not aﬀected by the magnetic ﬁeld, and then most of thegas in the disk accretes onto the NS equatorial region.
The HR would be low (for example HR < 0.09). According to increasing the luminosity, RA becomes smaller andthe size of Compton cloud would be smaller.
By attributing the diﬀerence of the size of Compton
cloud to the Alfvén radius, we can estimate the magnetic
ﬁelds of neutron star. Since HR of 4U 1820−30 and
4U 1735−44 is large, Bs is estimated as B ≥ 2.5×108 G.
The upper limit of B is given for other seven sources.
These results are consistent with the previous results.
4.

Transition between Horizontal Branch and Normal
Branch in Z sources
We investigated long-term variations in the hardnessluminosity diagram (HLD) of six Z sources by using
MAXI/GSC and Swift/BAT. We conﬁrm the secular motion of the Z-pattern through the HLD in four sources

106

r

1

(3)

R_c

P1

0.5
P2

(1)

0.2

R_A

Mdot [Eddington]

2

A
R_
(2)

The change of transition luminosity with the luminosity decreasing caused the shift of Z-pattern from Cyglike sources to Sco-like sources. The inner disk radius
becomes smaller as the transition luminosity decreases(Lin et al. 2009). Assuming that the large inner disk
radius (∼ 20 km) of Z sources represent the Alfvén radius of the neutron star with the surface magnetic ﬁeldof ∼ 109 G, the change of transition luminosity can be
attributed to a gradual change of accretion disk fromthe radiation-pressure-dominated disk to gas-pressuredominated disk. At the highest luminosity, the Z-patternshows Cyg-like with a long HB. At low luminosities, it
becomes Sco-like with a short HB. Figure 19 in Lin etal. (2009) shows the change of the inner disk radius at
the transition point of HB and NB. The radius increases
from Sco-like to Cyg-like. Here we consider the possibility of which the Alfvén radius determines the inner radius. We calculated Alfvén radius is for two special cases.
At high luminosity, the disk would be dominated by theradiation pressure [(1)R A line in ﬁgure 3]. On the other
hand, the lowest luminosity in Sco-like, the disk would be
gas pressure dominated [(2)R A line in ﬁgure 3]. Here,we consider the possibility of disk evaporation. Whenthe radiation pressure overcomes gravitational force, thedisk evaporation would occur. So, we also calculated theevaporation radius [(3)R cr line in ﬁgure 3]. When thesource is bright (Cyg-like), the inner radius moves along(1)R A line. This is the HB of Cyg-like. When the massaccretion rate reaches the evaporation line of (3)R cr,
evaporation starts (point P1). After then, the inner radius moves along (3)R cr. This is the NB of Cyg-like.
In a bright Cyg-like, the inner radius is large and theevaporation point is at high luminosity. As a result, the
HB is long. In a low luminosity as Sco-like, the disk becomes gas pressure dominate and the Alfvén radius (i.e.
inner disk radius) is shown as (2)R A line. In this case,the HB is short and the transition occurs at lower lu-

ISCO(Sco X-1, GX 17+2, Cyg X-2, and GX 5−1). Especially,
location of the HB and the NB shifted in the time scale ofseveral hundred days. Also, a transition luminosity from
HB to NB changed with the shift. We propose a simple
model of transition between HB and NB. Since the luminosity is close to the Eddington luminosity, the innerpart of the accretion disk is considered to be a radiationpressure-dominated disk. When the luminosity is evenhigher, the radiative force overcomes the gravity on thesurface of the accretion disk, and the surface starts toevaporate. We interpret that the diﬀerence of HB and
NB comes from the existence of the evaporation. In the
HB, the evaporation does not occur yet, because the luminosity is comparably low. When the source reaches atransition luminosity from HB to NB with the luminosity increasing, the disk evaporation will start. Namely,
the evaporation will occur in NB.

0.1
10

20
Radius [km]50

Fig. 3. The Alfvén radius (1) and (2) for B = 109 G (Campana etal. 1998), and the critical radius (3) (Fukue 2004) as a functionof the accretion rate. (1) is for radiation pressure dominated disk[equation (18) in Campana et al. 1998] and (2) for gas pressure
dominated disk [equation (14) in Campana et al. 1998].

minosity (point P2). In summary, we suggest that thediﬀerence of HB and NB might be explained by the disk
evaporation.
References
Asai, K., Mihara, T., Matsuoka, M., & Sugizaki, M.
2015, PASJ, 67, 92
Asai, K., Mihara, T., Mastuoka, M., & Sugizaki, M.
2016, PASJ, 68, 50
Barthelmy, S. D., et al. 2005, Space Sci. Rev, 120, 143
Campana, S., Colpi, M., Mereghetti, S., Stella, L., &Tavani, M. 1998, A&AR, 8, 279
Christian, D. J., & Swank, J. H. 1997, ApJs, 109, 177
Fukue, J. 2004, PASJ, 56, 569
Galloway, D. K., Muno, M. P., Hartman, J. M., Psaltis,
D., & Chakrabarty, D. 2008, ApJs, 179, 360
Gehrels, N., et al. 2004, ApJ, 611, 1005
Hasinger, G. & van der Klis, M. 1989, A&A, 225, 79
Kirsch, M. G.,et al. 2005, Proc. SPIE, 5898, 22
Kuulkers, E., & van der Klis, M. 2000, A&A, 356, L45
Lin, D., Remillard, R. A., & Homan, J. 2009, ApJ, 696,
1257
Lin, D., Remillard, R. A., Homan, J., & Barret, D. 2012,
ApJ, 756, 34
Liu, Q. Z., van Paradijs, J., & van den Heuvel, E. P. J.
2007, A&A, 469, 807
Matsuoka, M., et al. 2009, PASJ, 61, 999
Mihara, T., et al. 2011, PASJ, 63, S623
Mineshige, S., & Osaki, Y. 1983, PASJ, 35, 377
Mineshige, S., & Osaki, Y. 1985, PASJ, 37, 1
Smak, J. 1982, Acta Astronomica, 32, 213
Sugizaki, M., et al. 2011, PASJ, 63, S635
Tuchman, Y., Mineshige, S., & Wheeler, J. C. 1990, ApJ,
359, 164

107

Overview of Spectral Change in NS-LMXB
Hiromitsu Takahashi
Department of Physical Science, Hiroshima University, Higashi-Hiroshima, Hiroshima, Japan
E-mail (HT): hirotaka@astro.hiroshima-u.ac.jp
ABSTRACT
Neutron Star (NS)-Low Mass X-ray Binaries (LMXBs) are historically divided into two classes called
Z and atoll sources, according to the shape of their color-color diagrams. Z sources (including Sco X-1,
Cyg X-2, etc.) are always very bright close to Eddington luminosity, and show three distinct states
named as Horizontal, Normal and Flaring branches (HB, NB and FB). Some of atoll sources aretransient and sometimes switch the states between low/hard and high/soft states. After the discoveryof the bright transient NS-LMXB XTE J1701462, which followed the spectral states of both Z andatoll sources, it is confirmed that these two types of NS-LMXBs are determined mainly by the massaccretion rate. However, XTE J1701462 repeatedly showed all three Z states many times during thegradual outburst decay, and the physical origins are still unclear how Z sources change the threebranches of HB, NB and FB (almost) independent from the mass accretion rate. We present the recent
observational results of the high/soft state of the atoll source 4U 1608522, and explain the detailbehaviors of physical parameters, which may relate to the three states of Z sources.
KEY WORDS: binaries: general — stars: neutron — X-rays: starsSpectral States of NS-LMXBs

Neutron Star (NS)-Low Mass X-ray Binaries(LMXBs) are bright and they are observed fromthe beginning of X-ray astronomy. To classify their
behavior, color-color diagrams (CCDs) are used
and they are divided into two classes called Z andatoll sources (Hasinger et al. 1989). Z sources havea luminosity close to the Eddington limit (~1038
erg s-1), while atoll ones are transients fromquiescent to  Eddington luminosity. As shownin Figure 1, with increasing the accretion rate ,atoll sources show low/hard and high/soft states,
which are historically called as island and banana
states, respectively. Z sources are named as theyshow three states called Horizontal, Normal and
Flaring branches (HB, NB and FB). However, thedifferences of the three Z states are physically notunderstood yet. Among Z sources, there are two
sub classes called Cyg-like and Sco-like, since Cyg
X-2 and Sco X-1 are their representatives. Cyg-likeZ sources take “Z” shape in the CCDs, while Scolike ones show “”.
Since the recent bright transient NS-LMXBs(e.g., XTE J1701422 and MAXI J0556332)
showed both Z (Cyg-like and Sco-like) and atoll

behaviors during their outbursts, it is thoughtthat atoll sources become Z ones when  becomes
higher (Lin et al. 2009; Sugizaki et al. 2013). Thesetransients also repeated the three Z states manytimes during the gradual outburst decay. It
suggests that the three Z states are (almost)independent from  , although many previous
researches assumed that the Z behavior was
mainly controlled by the mass accretion rate.

Figure 1. States of NS-LMXB compared withthose of black hole binaries. This paper isabout the high/soft state shown in grey.

108

Compared with black-hole binaries (Figure 1),
which have very high or slim-disk states according
to  over the Eddington limit, it is unclear for
NS-LMXBs how physically different each of HB,
NB and FB is and what causes the transitions
among them at the (almost) constant .
In this paper, we summarize observational
results of the high/soft state of NS-LMXB andpropose the physical interpretations.


4U 1608 522 (atoll source): Outflow anddisk fluctuation

The emission of NS-LMXB in the high/soft state isnormally well represented by the combination of
the emission from the accretion disk and NS
surface (or boundary layer). The former is sum of
the black body emission from different
temperatures at different radii (= the multi-colordisk). The latter can be approximated by a singletemperature black body (Mitsuda et al. 1984;
Makishima et al. 1989).
We also used this “diskbb+bb” model in xspec inthe spectral analysis of the high/soft (upper
banana) state of an atoll source 4U 1608522
observed by Rossi X-ray Timing Explorer (RXTE)
satellite (Takahashi et al. 2011). Assuming theblack body emission, we obtain the luminosity andtemperature, and estimate the emission radius. As
a result, we realized the luminosity of the NS
surface (LNS) relatively decreases as that of theaccretion disk (Ldisk) increases (i.e., LNS/Ldisk
decreases from ~0.6 to 0.4 as the total luminosity
increases from ~1×1037 to 4×1037 erg s-1). Thisimplies that the ~20% fraction of the accretion
matter does not accrete onto the NS surface. Sucha behavior can be explained by the existence of an
outflow, which is probably caused by the radiationpressure of LNS + Ldisk (Figure 2).
When looking at deviations of the physicalparameters after subtracting the average trend,
we found two ways of fluctuations independent
from . One is that Ldisk (in the soft energy band)decreases and LNS (in the hard energy band)
increases simultaneously (Variable Luminosity
Branch: VLB). This VLB behavior makes the large
movement in the upper banana branch of the CCD.
The other is that there are no luminosity changes
but only the disk temperature (Tin) and radius (rin)fluctuate each other to keep Ldisk at constant(Constant Luminosity Branch: CLB).

Figure 2. The schematic view of the upper
banana state of atoll sources. Some fraction ofmatter from the accretion disk outflows anddoes not accrete onto the NS. The emission
radius of the NS surface is less than 10 km,
and may come from the equatorial region.

Figure 3. The same as Figure 2 but for thefluctuations independent from the massaccretion rate. Two ways are observed andtheir luminosities are variable (VLB) or
constant (CLB). At VLB, Ldisk decreases and
LNS increases simultaneously. At CLB, Tin andrin fluctuate with keeping Ldisk (and LNS) atconstant.
The former VLB suggests some fraction of thedisk emission is emitted not at the disk but at theNS surface, as shown in Figure 3. Such an
instability might occurred
due to photon
trapping at the inner edge of the disk (Ohsuga et

109

al. 2002). Then, the matter accretes onto the NS
with including the trapped photons, which arelater released as the emission from the NS surface
additionally. The CLB behavior appears only forthe temperature and hence the radius of theaccretion disk. This can be explained by thechange of the hardening factor, which might becaused by the disk structure (e.g., geometrical
thickness).
Relation between atoll and Z sources

Considering the atoll sources become Z ones at
the higher accretion rate and the above VLB and
CLB are independent from the mass accretion rate
(after subtracting the average trend), we thinkthat the three states of Z sources, whichtransitions are also (almost) independent from ,
are related to the VLB and CLB of the atoll ones.
Actually, the movements and positions of the
VLB and FB (at least Sco-like) in the CCDs are thesame (i.e., the harder flux increases in both states),
it is thought that the inner part of the accretion
disk drops to the NS without emitting efficiently.
The photon-trapping effect can occur morestrongly with the higher surface density.
The remaining CLB is the candidate of NB (andHB) of Z sources, when  increases. Z sources in
NB change only the harder flux (from the NS
surface) and keep the softer flux (from the disk)rather stable. The stable disk emission is the same
behavior observed in CLB. If CLB changes thegeometrical thickness of the accretion disk to
explain the behavior of Tin and rin, NB might havethe thinner disk when it is close to HB and thethicker one when close to FB. The cross-section
surface of the disk is suffered from the radiationpressure, and the matter outflows rather than theaccretion, if the pressure is large enough. The
thicker disk (close to FB) is thought to receive
more pressure and the more matter outflows,
resulting the less emission from the NS surface,since the accreting matter to the NS deceases. On
the other hand, the NS surface emission can behigher with the thinner disk (close to HB).
About HB behavior of Z sources, there areseveral models (e.g., disk state changes fromstandard to slim disk), but it should locate theextension of the NB physical explanation.
Furthermore, we need to understand why blackhole binaries do not show such three transitions

observed in the Z sources.
References
Hasinger, G. et al. 1989, A&A, 225, 79
Lin, D. et al. 2009, ApJ, 696, 1257
Sugizaki, M. et al. 2013, PASJ, 65, 58
Mitsuda, K. et al. 1984, PASJ, 36, 741
Makishima, K. et al. 1989, PASJ, 41, 531
Takahashi, H. et al. 2011, ApJ, 738, 62
Ohsuga, K. et al. 2002, ApJ, 574, 315

110

111

Unveiling the spectral transition of Aql X-1 from the hard to soft state
Ko Ono,1 Kazuo Makishima,2 Soki Sakurai,1Zhongli Zhang,3 Kazutaka Yamaoka4 , 5 and Kazuhiro Nakazawa1
1

Department of Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, JapanMAXI Team, Global Research Cluster, The Institute of Physical and Chemical Research, 2-1 Hirosawa,
Wako, Saitama 351-0198, Japan
3
Shanghai Astronomical Observatory, Chinese Academy of Sciences, 200030 Shanghai, China
4
Division of Particle and Astrophysical Science, Graduate School of Science, Nagoya University, Furo-cho,
Chikusa-ku, Nagoya, Aichi 464-8602, Japan
5
Institute for Space-Earth Environmental Research (ISEE), Nagoya University,
Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8601, Japan
E-mail(KO): ono@phys.s.u-tokyo.ac.jp
2

Abstract
The recurrent X-ray transient Aquila X-1 was observed with Suzaku on 2011 October 21 when the source
is in a rising phase of an outburst. The unabsorbed 0.1–100 keV luminosity increased from 3.5 × 1037 erg
s−1 to 5.1 × 1037 erg s−1 , and the 0.8–10 keV XIS count rate increased by a factor of ∼ 3, while the 15–
60 keV HXD-PIN count rate decreased to about one third. The spectrum showed a clear transition fromthe soft state to the hard state, changing in shape from a cutoﬀ power-law like shape to a convex thermal
one. Throughout the transition, the spectrum was successfully explained by a common model, consistingof a multi-color disk blackbody from a standard accretion disk and a blackbody emission from the neutronstar surface Comptonized by a corona. The transition is characterized by continuous changes of the model
parameters; the inner disk radius decreased from 31 km to 18 km, and the coronal temperature decreased
from 10 keV to 3 keV accompanied by a shrink of the corona toward the neutron star surface. Furthermore,
the fractional 0.1–100 keV luminosity carried by the disk emission increased from 21% to 56%.
Key words: accretion, neutron star, X-ray binary

1.

Introduction

Some Neutron-Star (NS) Low-Mass X-ray Binaries(LMXB), consisting of weakly magnetized NSs and lowmass stars, are known to undergo outbursts. When theluminosity exceeds ∼ 4% of the Eddington luminosity,
the spectrum changes drastically from the hard state exhibiting a power-law like shape with a cutoﬀ at > 10 keV
to the soft state showing a convex thermal one (Asai etal. 2012). Thanks to broad-band Suzaku data, the spectra in the two states have been described in a uniﬁed way
with three elements, multi-color disk (MCD) emission
from an accretion disk, a blackbody (BB) emission fromthe NS surface, and a hot corona that Comptonizes theBB photons weakly in the soft state and strongly in thehard state (Sakurai et al. 2012; Sakurai et al. 2014; Zhang
et al. 2014; Ono et al. 2016).
Since the interpretation of the soft state is rather established (Mitsuda et al. 1984) unlike that of the hardstate, to study how the hard state connects to the soft
state will reinforce the uniﬁed view. The best way forthis purpose is to resolve a spectral transition between
the two states, and conﬁrm that the spectral parameters in one state connect smoothly and continuously tothose in the other state. Although similar attempts havebeen made (Lin et al. 2007; Church et al. 2014), results
remained inconclusive due to the lack of low energy information needed to identify the MCD contribution.
The Suzaku observatory, which has wide band sensitivity covering both the MCD and the Comptonized BB
emission, is ideal for the above purpose. Among theSuzaku data of LMXBs, those of Aquila X-1 (hereafter
Aql X-1) acquired on 2011 October 21, in particular,
have turned out to be the most suited for our purpose,
because they caught a hard-to-soft spectral transition ina rising phase of an outburst (Ono et al. 2017). This
source is associated with an optical counterpart with amagnitude of V = 15 − 19 (Gottwald et al. 1991). The
distance has been constrained to be 4.4–5.9 kpc by its
type I bursts assuming an NS mass as MNS = 1.4 M
and the Eddington luminosity at the burst peak (Jonker

112& Nelemans 2004). In the present paper, the distance isassumed as 5.2 kpc.

2.

Data process

Fig. 1. Background subtracted light curves with 128 s binning of XIS
0+3 in 0.8–10 keV (top) and of dead-time corrected HXD-PIN in
15–60 keV (middle). The bottom panel shows the HXD-PIN vs
XIS hardness ratio. The time origin is 2011 October 21 12:51:33.
The arrows indicate type I bursts.

3. Analysis
As presented in ﬁg. 2, the obtained spectra show a clear
transition from P0 to P9. These 10 spectra were all ﬁtted with our uniﬁed model (§1), consisting of the MCD
plus Comptonized BB emission. XSPEC (version 12.9.0)
models diskbb, nthcomp (Zdziarski et al. 1996; Życki etal. 1999) and diskline (Fabian et al. 1989) were used.
The seed photons for nthcomp were chosen as BB expected to arise from the NS surface. The constructed
model, diskbb+nthcomp+diskline, was then multipliedwith an absorption factor wabs, with its column density
ﬁxed at 3.6 × 1021 cm−1 (Sakurai et al. 2014; Gatuzz et
al. 2016). The free parameters are the inner disk radius
Rin , the inner disk temperature Tin , the BB radius RBB ,the BB temperature TBB , the electron temperature Te ,the electron optical depth τ , the line centroid energy, theemissivity parameter, and the line normalization. The
inner disk radius of diskline was tied to Rin .
All the 10 spectra were reproduced successfully by themodel with χ2ν < 1.2 for 298 − 333 d.o.f. (ﬁg. 2). As
shown in ﬁg. 2, the 1–3 keV signal tripled from P0 to P9,
which is explained by an increasing MCD ﬂux. The spectrum above ∼ 10 keV, in contrast, decreased to about
one third, reﬂecting a decrease in the cutoﬀ energy, andhence in kTe .
The obtained parameters are summarized in ﬁg. 3.
The inner disk radius was derived using MCD normalization NMCD as
2

Rin = ξκ (NMCD )We used an archival Suzaku data set of Aql X-1 (ObsID 406010020) acquired from 2011 October 12:51:33 UT
for a gross duration of 79.9 ks. The XIS was operated
in the 1/4 window mode with a read-out time of 2.0 s.
The burst mode option was also employed to accumulate
photons for 0.5 s during the 2.0 s read-out time. Details
are given in Ono et al. (2017).
XIS 0 and XIS 3 events of GRADE 0,2,3,4 and 6 wereused. On-source events were accumulated in a 2 .5 circle
around the image centroid, but excluding events in the
innermost 1 circle to reduce the pile-up eﬀects (Yamada
et al. 2012). Background events were taken from an annular region with the inner radius of 4 and the outerradius of 5 . As shown in ﬁg. 3, we extracted spectrafrom 10 periods (P0,P1, · · · P10), but excluding type I
bursts.
HXD cleaned events were used as an on-source data after dead time correction. Non X-ray background (NXB)
events distributed by the HXD team was subtracted(Fukazawa et al. 2009) from the on-source data. CosmicX-ray background (CXB) was accounted for by adding
a ﬁxed CXB model to a spectral model. The data weredivided in the same way as the XIS.

1
2



D
10 kpc1(cos i)− 2

where D = 5.2 kpc is the distance, i = 45◦ is an assumed
inclination angle (Sakurai et al. 2012), and ξ = 0.412 andκ = 1.7 are the correction factor of the inner boundary
condition and the color-hardening of the disk, respectively (Kubota et al. 1998). The BB radius is calculated
assuming photon number conservation in the Comptonization process. In ﬁg. 3a, RBB < RNS < Rin is realized throughout, where RNS = 12 km is the NS radius,
implying a physically reasonable geometry. Through thetransition, the three temperatures became closer to one
another (ﬁg. 3b). The parameter values in the hardstate and those in the soft state are consistent with thetypical values found in previous observations of the respective states (e.g. Sakurai et al. 2012, 2014). Most importantly, all the parameters changed continuously andmonotonically. Thus, the hard state is smoothly connected to the typical soft state, with reasonable values.
Figure 3c shows the luminosities of the individual
physical processes. The disk luminosity Ldisk was calculated using 0.1–10 keV energy ﬂux of diskbb without
inclination correction. The BB luminosity LBB was obtained from the nthcomp photon ﬂux in 0.1–50 keV and

113

Fig. 3. Evolutions of the spectral parameters obtained from the best
ﬁt models, and the luminosities described in §3. Triangles, rectangles and circles represent the corona, BB and the disk, respectively.
(a) The radii. (b) The temperatures. (c) The total luminosity(stars), with ordinate on the left. The luminosity fractions of thethree processes are also shown with ordinate on the right.

Fig. 2. The Suzaku spectra of Aql X-1 in νF ν form from P0,
P6 and P9, acquired with XIS 0+3 (0.8–10 keV), HXD-PIN(15–60 keV) and HXD-GSO (>50 keV). They are ﬁtted withwabs*(diskbb+nthcomp+diskline), and the best ﬁt model is also
shown together with its components.

its BB temperature, assuming a photon number conservation and a spherical emission. The Comptonization
luminosity Lcomp was derived by subtracting LBB fromthe luminosity of the nthcomp at 0.1–100 keV. As thetotal luminosity Lt increased in ﬁg. 3c, Lcomp decreasedand Ldisk increased with a negative correlation. The XISvs HXD-PIN anti-correlation, ﬁrst revealed in ﬁg. 1, approximately reﬂects this Ldisk vs Lcomp anti-correlation.
The fraction of LBB did not change throughout and is
not much diﬀerent from 50%, which is predicted by thevirial theorem (§4.2).
4. Discussion
During the observation of Aql X-1 in its outburst, thetransition took place in ∼ 30 ks involving a drastic
spectral change. Throughout the subdivided 10 periods,
the spectra were successfully explained with a common
model of wabs*(diskbb+nthcomp+diskline). The obtained parameters evolved continuously and monotoni-cally from the values typical of the hard state to thoseof the soft state. As a result, the spectral components in
the hard state have been ﬁrmly identiﬁed with those in
the soft state.
4.1. Interpretation of the model evolution
As the luminosity increased, the inner disk edge approached the NS surface from Rin = 31 km to 18 km.
Consequently, RBB decreased from 11 km to 7 km. Thissuggests that a coronal accretion ﬂow in the hard stateis onto the whole NS surface, while that in the soft state
become more conﬁned to the NS equator. In association
with the above geometrical change, Te decreased due toweaker heating and stronger cooling of the corona. To
summarize, the state transition can be interpreted asa transition of the emission from that dominated by aComptonizing corona to that dominated by an opticallythick accretion disk.
4.2. Coronal accretion
In addition to the MCD and BB emission, we can further
investigate coronal accretion utilizing the coronal opticaldepth τ along the line of sight. Here, τ is expressed asτ = σtne dx(1)where σt is the Thomson scattering cross section andne (R) is the electron number density at a radius R. The

114

mass accretion rate can be expressed asṀ = 1.2S(R)Vr (R)μmp ne (R)(2)

where 1.2 is the average molecular weight, S(R) =4πR2 ζ is the cross section of the corona at R, 0 < ζ < 1
is a form factor, Vr (R) is the radial velocity at R and mp
is the proton mass. We assumed a spherical 
accretion forsimplicity. Using the free-fall velocity Vﬀ = 2GMNS /R
where G is the gravitational constant, and a dimensionless parameter g (0 < g < 1), we may write as
Vr = gVﬀ (R). Substituting Lt = GMNS Ṁ /RNS to eq.
(2), ne in eq. (1) can be expressed as ne = ne (g, R, Lt ).
Integrating eq. (1) from RNS to Rin ,
0.16
gζ =
τ



Rin
−1
RNS

Lt ( erg s−1 )
1037

(3)

is obtained. Substituting the obtained parameter valuesof τ , Rin and Lt , we derived gζ as in ﬁg. 4. The smaller
values of gζ in the soft state indicate that the scale heightof the corona become lower (smaller ζ), and/or the radial
coronal velocity is lower (smaller g). In any case, theradial velocity is much smaller than the free fall velocity
over the observation, assuming ζ is close to the unity.
Therefore, the corona is in fact close to the Keplerianmotion.
In order to compare the present results with those froma soft to hard state transition, gζ from Sakurai et al.
(2015) is also shown in ﬁg. 4. Similarly, i) gζ is as small
as ∼ 0.01 both in the soft and the hard state, and ii) gζ
is smaller in the soft state than in the hard state, but iii)
the change of gζ took place at a lower luminosity below∼ 2 × 1037 erg s−1 due to the hysteresis. In ﬁg. 4, thesource evolve clockwise along the small arrows when thesource undergoes an outburst, from the hard state to thesoft state, and from the soft state to the hard state.
References
Asai, K., Matsuoka, et al. 2012, PASJ, 64,
Fabian, A. C. et al. 1989, MNRAS, 238, 729
Church, M. J. et al. 2014, MNRAS, 438, 2784
Fukazawa, Y., et al. 2009, PASJ, 61, 17
Gatuzz, E., & Garcı́a, J. A. et al. 2016, A&A, 588, A111
Gottwald, M. & Steinle, H. et al. 1991, A&AS, 89, 367
Jonker, P. G., & Nelemans, G. 2004, MNRAS, 354, 355
Kubota, A., Tanaka, Y. et al. 1998, PASJ, 50, 667
Lin, D., Remillard, R. A. et al. 2007, ApJ, 667, 1073
Mitsuda, K., et al. 1984, PASJ, 36, 741
Ono, K., & Makishima, K. et al. 2016, PASJ, 68, S14
Ono, K., 2017, submitted for publication PASJ
Sakurai, S., Takahashi, H. et al. 2012, PASJ, 64, 72
Sakurai, S., et al. 2014, PASJ, 66, 10
Sakurai, S., 2015, PhD Thesis, The University of Tokyo

Fig. 4. The evolution of gζ, from the spectral ﬁt, as a function of
Lt . Filled circles are the present results and the open circles arethose taken from Sakurai et al. (2015).

Yamada, S., Uchiyama, H., et al. 2012, PASJ, 64, 53
Zhang, Z., & Makishima, K. et al. 2014, PASJ, 66, 120
Zdziarski, A. A. et al. 1996, MNRAS, 283, 193
Życki, P. T., & Done, C. et al. 1999, MNRAS, 309, 561

115

The MAXI/GSC X-ray Burst Catalog
Takanori Sakamoto,1 Hikaru Ohtsuki,1 Motoko Serino,2Atsumasa Yoshida1 and MAXI team
1

Aoyama Gakuin University, 5-10-1 Fuchinobe, Chuo-ku, Sagamihara-shi Kanagawa 252-5258, Japan
2
RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(TS): tsakamoto@phys.aoyama.ac.jp
Abstract

We present the X-ray burst (XRB) catalog based on the MAXI/GSC observations. Between August
15, 2009 and November 30, 2015 (six years and three months), we identiﬁed 329 XRBs from 24 sources.
The majorities of XRBs were detected from H 1636−536 (79 XRBs) and GS 1826−238 (139 XRBs). We
compared burst durations, radiation energies and burst frequencies of those XRBs with the persistent
ﬂuxes near the bursts, and found relationships among them. We discuss the general picture of XRBs
based on the MAXI/GSC observations.
Key words: workshop: MAXI/GSC — X-ray bursts — Catalog

2. Light curve analysis
First, we used the GSC light curve data generated by
the MAXI on-demand process1 for the scan bases. Using
this light curve, we extract the candidate times for the
XRB events. We extracted 1765 XRB candidates from
24 sources in this stage 1 selection (Figure 1).
Second, the 1-s light curve is generated for the above
XRB candidates from the GSC event data. The background has been subtracted and the eﬀective area is corrected for the generated light curve. Then, the persistentemission component is subtracted from it to generate thelight curve only containing an XRB emission.
*1 http://maxi.riken.jp/mxondem/

H1636-536
3

On-demand light curve
X-ray burst candidates
Estimated persistent emissions

2.5

2Rate [photons/cm2/s]

1. Introduction
Most of the low mass X-ray binaries (LMXB) which are
the source of X-ray bursts (XRB) show a steady emission from the accretion disk. Various characteristics of
XRBs depend on the environment of the neutron starsurface and the composition of the fuel. The composition depends on the mass accretion rate ṁ. The massaccretion rate is known to be proportional to the brightness of the persistent emission (Galloway et al. 2008).
However, the rule of the magnetic ﬁeld and the eﬀect ofa non-uniform mass accretion to a neutron star surface
are still not clear. Furthermore, the relation between apersistent emission and XRBs is still unresolved problem. In this paper, we investigate the relation between
a persistent emission and XRBs using the MAXI/GSC
data.

1.5

1

0.5

0

-0.5
55000

55500

56000

56500

57000

MJD [day]Fig. 1. Example of the stage 1 selection of the burst candidates ofH1636−536. Red is the scan-based GSC light curve generated by
the MAXI/GSC on-demand process. Green marks are the candidates of XRBs. Purple marks shows the persistent emission level.

Finally, the light curve is ﬁtted by a fast rise and exponential decay model (Figure 2). Based on the bestﬁt parameters, we extract the following quantities: t90
duration, 2-10 keV ﬂuence, a peak photon ﬂux Fpk , abolometric ﬂuence in 0.2-200 keV Eb , the characteristic
time scale τ (=Eb /Fpk ), the normalized burst ﬂuence
Ub (=Eb /FEdd ) where FEdd is the ﬂux corresponding
to the Eddington luminosity and γ (=Rper /REdd ) where
REdd is the photon rate corresponding to the Eddington
luminosity. τ , Ub and γ are the parameters deﬁned in
Galloway et al. (2008). We also calculate the burst rate
V which is the XRB burst rate normalized by the total

116

number of the MAXI/GSC scans.

Gamma vs Ub

10

6
4

Ub = Eb / FEdd
, BN= 4.873

SNR
2
0

ST= í12.13, PT= í11.74, DT= 4.608

6
4
2
0

RATE [photons/s/cm2]

H1636í536 light curve bin=1s with SNR

í20
0
20
Time since 2009í08í30T07:16:49.3 [s]H1636-536
H1735-444
4U0513-40
4U1608-52
AqlX-1
Terzan2HETEJ1900.1-2455
NGC6624
4U1916-053
M15
4U1705-440
RXJ1832-33
GS1826-238

1

0.01

0.1

Normalized persistent flux gamma

Fig. 4. Relationship between the normalized burst ﬂuence Ub andthe normalized persistent ﬂux γ.

Gamma vs Burst rate V

Fig. 2. The 1-s light curve of the XRB from H1636−536 createdfrom the MAXI/GSC event data (red). The light curve is ﬁtted
by a “burst” model (solid black line).

5
4

3.

Relationship between the persistent emission and thecharacteristics of XRBs
Figure 3 shows the relationship between the characteristics time scale τ and the normalized persistent ﬂux γ.
As can be seen in the ﬁgure, XRBs which have longerτ tend to have larger persistent ﬂux γ. Whereas, XRBs
with shorter τ tend to have smaller persistent ﬂux γ.

Tau [s]Gamma vs Tau
H1636-536
H1735-444
4U0513-40
4U1608-52
AqlX-1
Terzan2HETEJ1900.1-2455
NGC6624
4U1916-053
M15
4U1705-440
RXJ1832-33
GS1826-238

10

1

0.01

0.1

Normalized persistent flux gamma

Fig. 3. Relationship between the characteristic XRB time scale τ andthe normalized persistent ﬂux γ.

The relationship between the normalized burst ﬂuence
Ub and the normalized persistent ﬂux γ is shown in Figure 4. We found that XRBs with small persistent ﬂux
tend to release a large amount of energy as XRBs. Whenγ gets greater than 0.03, there is no diﬀerence in the ﬂuence Ub of XRBs.

Burst rate V

3
2

0.001

1
0

0.00010.01

0.1

Normalized persistetn flux gamma

Fig. 5. Relationship between the burst rate V and the normalized
persistent ﬂux γ.

Figure 5 shows the color map between burst rate V
and the normalized persistent ﬂux γ. The color shows
the number of sources. The highest burst rate is evident
for XRBs which show the middle range of the persistent
ﬂux level in the sample. Then, the burst rate decreases
in the order of the samples showing the lowest and thehighest persistent ﬂux.
4. Summary
Based on the results between the burst characteristics
and the persistent emission shown in Section 3., we believe the wide variety of the XRB characteristics could
be understood by the diﬀerence of nuclear reaction of thematters accreting to the neutron star surface. Further
continuation of the observations by MAXI is crucial to
conﬁrm our general view of XRBs.
References
Galloway, D.K. et al. 2008, ApJS, 179, 360

117

Hard X-ray study of low mass X-ray binary Sco X-1
with Suzaku
Chiho Okada1, Hiromitsu Takahashi1, Tatsuya Nakaoka1, Takao Kitaguchi1, Yasushi Fukazawa1,
Tadayasu Dotani2, Yoshitomo Maeda2, Satoru Katsuda2, Hideyuki Mori3, Murakami Hiroshi4
1 Hiroshima University
2ISAS
3UMBC
4Tohoku Gakuin University

E-mail (CO): okada@astro.hiroshima-u.ac.jp
ABSTRACT
Low-mass X-ray binaries with neutron stars, whose luminosity are close to the Eddington limit (1038erg/s), are called as Z sources. It is known that Z sources transit among three states (branches) on
their color-color diagrams (CCD); Horizontal branch (HB), Normal branch (NB), and Flaring branch(FB). Apart from the thermal emission, the hard X-ray tail above 50 keV from Sco X-1 has beenreported by RXTE and INTEGRAL satellites. However, detection of the hard emission relies on thetime variability and instrument calibration uncertainties, and thus the characteristics of the hard
tail are not understood well. We think that clarifying this phenomenon leads to understanding aphysical picture around the neutron star and their neighboring accretion disk. Since Sco X-1 is themost brightest object in the X-ray band, it is suitable to examine time variation of the hard tail withthree state transitions. Therefore, we performed an observation of Sco X-1 by Suzaku on March 9-12,
2015. We detected that Sco X-1 transited among three states (HB, NB, FB) of Z sources during theobservation by making the CCD of XIS and HXD-PIN data. We also found that inverse Compton
scattering of black-body seed photons from the neutron star surface can reproduce the observed
spectrum only less than 35 keV. The hard tail was needed to exist above 50 keV with the flux of ~4×10-10 erg/cm2/s in 20-200 keV at average.
KEY WORDS: X-ray binaries
1.

Introduction

From Sco X-1, emission of hard X-ray tail above
50 keV has been reported by RXTE, INTEGRAL
satellites. Many of these results show that thehard X-ray flux has decreasing trend from HB,
NB to FB (e.g., [1]). There is also a report thatsome RXTE observations of Sco X-1 showed thestate of top FB, which has the very hardcomponent of the power-law index of ~1.4 [2].
Therefore, we analyzed Suzaku data of Sco X-1,
and investigated the relationship between statetransition and hard tail.
2.

Observation and Result

Reproducibility of the PIN background model is
not good in 2015, because the backgroundobservation was limited due to the electricity
shortage of Suzaku. Therefore, we corrected thebackground model carefully to reproduce

earth-occultation data within 1% level, using GK
Per (previous observation) and Sco X-1. Then, thesystematic uncertainty of the background is
added in the following analysis, PIN : 3~4 % in
35-70 keV and GSO : 1% in 50~500 keV.
2.1. Spectral fitting of average spectrum
We first fit the average spectrum of HXD(PIN+GSO) data. The net exposure time is 30 ks
and the systematic error of 0.3% is added forfitting models.
The result shows the high energy region above
35 keV cannot be reproduced by only the single
inverse Compton emission (compTT model) with
2/d.o.f = 87.1/65, and requires the additional
power-law component (2/d.o.f = 65.52/63). The
null hypothesis not to exist the power-law
emission is 0.01% by F test. Figure 1 and Table 1
show the fitting results of compTT+power-law
model.

118

Figure 3. color-color diagram
Figure 1. fitting with compTT + power-law modelcompTT

power-lawTemperature (keV)(Seed photon)
Temperature (keV)(Scattering electron)
Optical depth

1.7+0.20.1

Flux (10-9 ergs/cm2/s)

10 1
3

Photon Index 

4+1
2
4+4
2

Flux (20-200 keV)
(10-10 ergs/cm2/s)

3.0 +0.2
0.2
3.9 +0.7
0.5

Table 1. best fit parameters with 90% errors

2.2. Time variability

2.3. Spectrum fittng of each state
We obtain the PIN spectra at each state andevery 96 minutes (= 1 satellite orbit), and fittedwith compTT + power-law model. Figure 4 shows
the obtained 20-200 keV power-law flux (red:each state, blue: 96 minutes from HB, NB to FB
shown in the green arrow of Figure 3.). This
result confirms that the flux of the hard taildecreases toward FB. The very hard emission
reported in top FB was not observed in the
PIN+GSO spectrum of FB (=1.4 fixed). The
upper limit of the 20-200 keV flux is 6 ×10-10
erg/cm2/s, which is a half of the value in [2].

We plot the light curve and CCD of Suzaku datawith 128-s binning for three energy ranges of 0-2,
4-8, and 15-20 keV, to examine the time
variability of the hard tail (Figure 2 and 3). Thisshows that Sco X-1 was highly variable during
Suzaku observation. In 15-20 keV, the averagerate is ~30 counts/s, but it reaches ~80 counts/s
at the maximum. CCD shows the three states of
HB, NB and FB in ““ shape.
Figure 4. flux of power-law model (=2.7 fixed)

3.

Discussion

We observed that Sco X-1 exhibited the threestates of Z sources during the Suzaku observation,and the flux of the hard tail changed anddecreased from HB, NB to FB. The hard X-ray
tail above 50 keV is interpreted as the inverse
Compton scattering of high energy electrons.
Such electron clouds seem to exist only in HB and
NB, and disappear in FB.
Figure 2. light curve

References[1] T. Di Salvo, et al. 2006, ApJ, 649, L91-L94[2] L. Titarchuk, et al. 2014, ApJ, 789, 98

119Black holes and Neutron stars in Low Mass X-ray binaries
Chris Done1
1

ISAS/JAXA and University of Durham, UK
E-mail(CD): chris.done@durham.ac.uk
Abstract

Black holes and neutron stars both show dramatic spectral and timing evolution as a function of massaccretion rate. This is seen most clearly in the transient outbursts of these systems, which can be followedwith MAXI and other all sky monitors. I will discuss the similarities and diﬀerences between the neutronstars and black holes, showing how to ﬁt this together in a picture of a very similar accretion ﬂow with avery diﬀerent boundary condition - an event horizon in black holes, versus a surface in neutron stars.

120

121

Understanding superbursts
Jean in ’t Zand
SRON Netherlands Institute for Space Research, Sorbonnelaan 2, 3584 CA Utrecht, the Netherlands
E-mail: jeanz@sron.nl
Abstract
Superbursts were discovered at the beginning of this millennium. Just like type-I X-ray bursts, they
are thought to be due to thermonuclear shell ﬂashes on neutron stars, only igniting much deeper. With
respect to type-I bursts, they last 103 times longer, are 103 as rare, ignite 103 times deeper (in column
depth) and are thought to be fueled by carbon instead of hydrogen and helium. Observationally, they aresometimes hard to distinguish from intermediate duration bursts which are due to pure helium ﬂashes oncold neutron stars. So far, 26 superbursts have been detected from 15 neutron stars in low-mass X-ray
binaries that also exhibit type-I bursts. They are very diﬃcult to catch and only 2 have been measured
with highly sensitive instrumentation. Superbursts are sensitive probes of the neutron star crust and theaccretion disk. The superburst phenomenon is not fully understood. Questions remain about the natureof the fuel, the collection of that fuel and the ignition conditions. The current state of aﬀairs is reviewed
and possible resolutions that lay ahead in the future discussed.
Key words: Stars: neutron – X-rays: bursts, binaries – Nuclear reactions

1. Introduction
Since 1969, ’type-I’ X-ray bursts are being detected fromspace-borne observatories between roughly 1 and 10 keV(Belian et al. 1972; Grindlay et al. 1976; Matsuoka et al.
1980; Makishima et al. 1981; Gottwald et al. 1986; Lewin
et al. 1993; Strohmayer & Bildsten 2006; Galloway et al.
2008). These bursts are due to thermonuclear shell
ﬂashes of hydrogen and helium in the freshly accreted
upper layers of neutron stars in low-mass X-ray binaries (LMXBs; Hansen & van Horn 1975; Lamb & Lamb
1978; Wallace & Woosley 1981; Fujimoto et al. 1981), likethose on white dwarfs are responsible for classical novae.
The durations of most type-I bursts are between a few
seconds and a few minutes. The time proﬁle of the X-ray
emission is typically a fast rise of duration a few seconds
and an exponential-like decay. The shell ﬂash occurs atlarge ( 1) optical but small linear depth (roughly 1m) below the photosphere and temperatures may riseto order 1 GK. By the time the heat wave reaches thephotosphere, the temperature is a few tens of MK andthe typical peak energy of the photon spectrum is 5-10
keV. The ﬂash itself lasts a fraction of second, although
it may take a few seconds to engulf the whole neutronstar. The burst decay phase essentially is the coolingof the neutron star and actually does not follow an exponential function but a power law (e.g., Cumming &Macbeth 2004; in ’t Zand et al. 2014). The duration of
the cooling is primarily set by the amount of mass thatis heated up or, in other words, the column depth of theignition (usually between 108 and 109 g cm−2 ).
In the fall of 1996, the just launched two Wide Field
Cameras (WFCs; Jager et al. 1997) onboard the ItalianDutch BeppoSAX observatory (Boella et al. 1997) conducted a nine-day long observation of the Galactic center region. Thanks to the 40 × 40 square degrees ﬁeldof view per camera, the WFCs could simultaneously observe about half the Galactic low-mass X-ray binary population. One of the LMXBs in the ﬁeld of view was
4U 1735-44. Figure 1 shows the nine-day light curve
resulting from that observation. It includes a remarkable feature on August 22, 1996: a fast rise exponentiallike decay phenomenon with a duration of a few hours.
Cornelisse et al. (2000) found this, and the spectrum,very reminiscent of a type-I burst except for the duration which is 103 as long. They proposed this as thelongest thermonuclear burst ever observed. This marks
the discovery of superbursts.
The discovery gave rise to searches in archival datafrom WFC and the All-Sky Monitor (ASM) on RXTE,
which yielded 12 more very long bursts that were published between 2001 and 2004 (see Table 1). Wijnands(2001) introduced the term ’superburst’. Cumming &Bildsten (2001) and Strohmayer & Brown (2002) introduced the ﬁrst explanation of the phenomenon as a thermonuclear shell ﬂash fueled by carbon at a column depth∼ 103 times deeper than for type-I bursts.

122

Fig. 1. Nine-day long light curve of 4U 1735-44 as observed with BeppoSAX-WFC, with the ﬁrst detection of a superburst on August22, 1996 (from Cornelisse et al. 2000)

In this paper, we brieﬂy discuss the observational facts(§2.) and theoretical considerations (§3.) that deﬁne ourcurrent understanding of superbursts, summarize the importance of superburst research (§4.) and touch on future
prospects (§5.).

2.

Observational overview

2.1. Catalog & recurrence time
Table 1 presents a list of all 26 superbursts and theircharacteristics that have been reported up to January
2017. The most recent one is reported in these proceedings (Iwakiri et al. 2016a, see also Iwakiri et al. 2016b). It
is from 4U 1705-44, a LMXB that was already predicted
to be a superburster a decade ago (in ’t Zand et al. 2003).
The 26 superbursts are emitted by 15 low-mass X-ray binaries that are also emitters of type-I bursts. Currently,
MAXI on the ISS is the most eﬃcient superburst discovering machine. All six superbursts since 2011 werediscovered with this device.
Five sources have exhibited multiple superbursts.
They have high accretion rates except 4U 0614+09. The
range of superburst recurrence times is between 10 d(for GX 17+2) to 10.5 yr (for 4U 1820-30). The average recurrence time is 4 yr, but this should be taken
as an upper limit because there are data gaps. in ’t
Zand et al. (2003) perform a statistical analysis of therecurrence time and ﬁnd 2 ± 1 yr. Keek et al. (2006)
determine on a source-by-source basis and on the basisof the BeppoSAX-WFC database a lower limit to therecurrence time of usually 2 months. These determinations can probably be improved upon considerably withthe much larger data sets that is now available throughfor instance INTEGRAL and MAXI.

2.2. Host binaries
While in the early years (2000-2004) host binaries ofsuperbursts were all found to be LMXBs that are persistently accreting at a level of at least 0.1 times the
Eddington limit, the picture changed as the data grew.
There are now 4 transients among the 15 superbursters,
and one system (SAX J1828.5-1037) with an unknown
nature, but for certain with a low long-term average accretion rate like the 4 transients. Of the 7 host binaries
with (tentatively) known orbital periods, 2 are ultracompact X-ray binaries (UCXBs), meaning that the composition of the donors, and therefore that on on the neutronstar, is strongly deﬁcient in hydrogen (4U 1820-30 and
4U 0614+09).
There are a number of proliﬁc bursters that are semipersistent and did not exhibit a superburst yet: 4U 172834, EXO 0748-676 (oﬀ since 2011), 4U 1702-429, 1E
1724.3045 (in Terzan 2), A 1742-294, 4U 1812-12, GS
1826-24 and Cyg X-2. While it is too early to derivea physical meaning of this, it is something to keep inmind. It may be related to the question of fuel accumulation for superbursts and ignition conditions. in ’t
Zand et al. (2003) did an investigation of the average
α parameter among a number of persistent bursters. α
is the ratio of the gravitational energy released by accretion since the last burst to the nuclear energy released through the present burst. It should be between
about 30 and 200 (Lewin et al. 1993). in ’t Zand et al.
(2003) found a clear distinction between superbursters
and non-superbursters. The former ones have a signiﬁcantly higher α value (>∼ 1000).
2.3. Distinguishing superbursts from other long bursts
After the discovery of superbursts, another kind of thermonuclear X-ray burst was discovered that is also long
but generally not as long as superbursts: intermediate
duration bursts (in ’t Zand et al. 2005; Cumming et al.
2006), the qualiﬁcation intermediate referring to a duration between that of type-I bursts and superbursts (seeFig. 2). Since this may incur uncertainty in the identiﬁcation of short superbursts, we discuss this somewhat
more.
Most intermediate duration bursts are thought to result from the ignition of thick helium piles on cold neutron stars. Due to the low temperature, ignition isreached at higher pressure and larger column depths,
and the helium needs much more time to reach ignition(days to weeks instead of hours) and bursts last longer.
The thicker piles and the fast 3α helium-burning nuclear
reaction will usually result in very high nuclear powersthat easily surpass the Eddington limit and result in verystrong photospheric expansion. The low temperatures go
hand in hand with low accretion rates. All these circumstances are found in ultracompact X-ray binaries withorbital periods shorter than about one and a half hour

?
y

1137

2005-03-12(ASM)
2014-11-03(MAXI)
1999-01-09(WFC)
2005-05-05(ASM+HETE)
1996-06-19(ASM)
1997-07-13(ASM)
1999-05-26(ASM)
2001-02-22(PCA+ASM)
2016-10-22(MAXI)
1996-09-23(WFC)
1996-08-22(WFC)
1998-06-09(ASM)
1996-09-14(WFC)
1999-09-23(WFC)
1999-10-01(WFC)
2000-09-08(WFC)
2011-10-24(MAXI+BAT)
2011-02-13(JEMX+MAXI)
1999-09-09(PCA)2010-03-17(MAXI+ASM)∗
1997-02-28(WFC)
1999-08-09(ASM)
2008-10-14(ASM)
2011-12-06(MAXI)
2011-11-12(MAXI)
2013-07-20(MAXI)

Date (Instr.)

y
y

y
y

y

y

y
y

Onset
obs.?
2.1
5.2
6.0
4.5
3.1
1.8
2
1.5
2.2
2.7
1.4
1.6
1.9
1.0
0.7
2.2
10
4.2
1
0.5
1.2
3.6
1.4
2.3
2.3
4.3

Decay
Time
(hr)
< 0.01

> 0.1
2.8
0.4
0.5
1.0
> 0.9
0.8
0.7
> 1.0
1.4
1.5
0.8
1.0
1.3
1.7
1.8
0.7
3
3.4
> 3.3
1.6
> 0.19> 0.26
0.9
0.7
1.0
0.1
0.25
0.2
1
1
1
1
< 0.01
0.1
0.1
0.15
0.2
0.15
0.13
0.21
< 0.01
0.1

0.13
0.03

Ṁ
(edd.)

Lpeak(1038
erg s−1 )+389-711/+25-168/+167+1549-162/+34+309+55-6/+34../+374-62/+94
+2
+10
+2
+12

-51/+125-57/+104+96-122/+68-27/+15-2/+23-367/+19

Nearestburst (d)

2.3

4

2.1

10

>1

1.8

1.9
2.6

2.6

1.5
1.6

5

E
(1017
erg g−1 )

0.55

1

1.0

0.6

1.0
1.3

0.48

2.7
2.8

0.2

y
(1012
g cm−2 )1,a
2
3, 4, b
5
6, c
7
8
9, 4
10
11, 4
12, 4, d
13
14, e
14
14, 4
14, 4
15, 16
17, 2
18, f
19
20, 4
7
7
10
21
2, g

Ref.‡

According to Serino et al. (2016), this is not a superburst. It is listed here, because we believe that the alternative, it being an intermediate duration burst, isless likely given the high accretion rate. ‡ References (numeric for superburst data, alphabetical for orbital period): 1 - Kuulkers et al. (2010), 2- Serino et al.
(2016), 3 - in ’t Zand et al. (2003), 4 - Cumming & Macbeth (2004), 5 - Keek et al. (2008), 6 - Wijnands (2001), 7 - Kuulkers (2009), 8 - Kuulkers et al. (2004),
9 - Strohmayer & Markwardt (2002), 10 - Iwakiri et al. (2016a), 11 - Kuulkers et al. (2002), 12 - Cornelisse et al. (2000), 13 - Kuulkers (2002), 14 - in ’t Zandet al. (2004), 15 - Altamirano et al. (2012), 16 - Serino et al. (2012), 17 - Chenevez et al. (2011), 18 - Strohmayer & Brown (2002), 19 - in ’t Zand et al. (2011),
20 - Cornelisse et al. (2002), 21 - Asada et al. (2011), a - Shahbaz et al. (2008), b - Motch et al. (1987), c - van Paradijs et al. (1990), d - Corbet et al. (1986), e Bandyopadhyay et al. (2002), f - Priedhorsky et al. (1986), g - Chevalier & Ilovaisky (1991)∗

SAX J1828.5-1037
Aql X-1

Ser X-1

11yy

EXO 1745-248
SAX J1747.0-2853
4U 1820-30

279

10d?

∼

228

4U 1705-44
KS 1731-260
4U 1735-44
GX 3+1
GX 17+2

y236

4U 1254-69
4U 1608-52
4U 1636-53

Porb
(min)

51?

Transient?

4U 0614+09

SourceTable 1. Catalog of 26 superbursts from 15 sources, detected up to January 2017.

123

VXSHUEXUVWVLQWHUPHGLDWH
GXUDWLRQEXUVWVPL[HG++HEXUVWV Î+HEXUVWV Î

124

Fig. 2. Preliminary histogram of burst durations as determined frombursts in the MINBAR archive (Galloway et al., in prep.)(Nelson et al. 1986) and that is indeed what observations
show (in ’t Zand et al. 2007; in ’t Zand & Weinberg 2010).
Peng et al. (2007) predict a small regime in accretion
rate of H-rich systems (≈ 0.003 times Eddington) where
also long helium bursts can happen after long series ofpure H bursts. However, this is a very small range of allowed accretion rates and pure H-bursts have never been
detected yet. Nevertheless, there are sporadic reports ofintermediate duration bursts from H-rich systems (Degenaar et al. 2010; Chenevez et al. 2007).
Identifying superbursts, particularly when data coverage is sporadic, can be cumbersome. One should, for instance, be careful about long bursts that are discoveredfrom UCXBs at low mass accretion rates. These may
be intermediate duration bursts. In fact, in a few cases
superburst detections had to be revised, see the careful
evaluation by Serino et al. (2016) . Furthermore, in myopinion this qualiﬁes the identiﬁcations as superbursts
in 4U 0614+09 and SAX J1818-1036 as less certain (seealso Kuulkers et al. 2010 for 4U 0614+09).
2.4. Peak luminosities
Roughly 20% of type-I bursts have peak luminosities
near the Eddington limit (Galloway et al. 2008): 2.0 ×1038 erg s−1 for H-rich atmospheres and 3.4×1038 erg s−1
for H-poor atmospheres. The situation is diﬀerent forsuperbursts. All superbursts except the PCA one from
4U 1820-30 (Strohmayer & Brown 2002) and possible
the JEM-X burst of SAX J1747.0-2853 (Chenevez et al.
2011) are sub-Eddington. This immediately shows thatthe fuel layer is not burning completely.
2.5. Precursors
Most superbursts are discovered with low duty-cycle instruments, particularly the ASM on RXTE and MAXI

Fig. 3. PCA-measured light curves of two bursts from 4U 1820-30.
Top panel: ﬁrst 30 s of the superburst. Bottom panel: a type-I
burst from the same source. From Strohmayer & Brown (2002).

on the ISS. These devices observe more than 80% of thesky every 90-min satellite orbit, but only for about 1 min.
Therefore, it is easy to detect superbursts because theygenerally last longer than 90-min, but it is also diﬃcult
to catch the onset of superbursts. The onset has been
observed in 8 of the 26 superbursts. Interestingly, in eachof these cases the onset is marked by a short burst. This
is often called a precursor, but actually there is only one
case (the PCA superburst of 4U 1820-30) where there istruly a brief period without burst emission between the
precursor and the main burst.
Figure 3 shows the onsets of a superburst and an ordinary burst from 4U 1820-30, detected with the highthroughput PCA. The superburst onset is characterized
by the precursor (from 5-18 s) and the superburst (starting at 18 s), and the dips in both these bursts. The dip
in the precursor is very short (less than the time resolution of 81 s) and possibly not complete to the pre-burst
ﬂux level. The dip in the superburst drops to below thepre-burst level. Both of these dips are consistent withphotospheric expansion (Lewin et al. 1984) with adiabatic cooling, whereby the cooling in the second dip is
so strong that the X-ray signal is lost (Keek 2012). The
drop to below the pre-burst level is due to the photosphere covering up the X-ray emitting part of the accretion disk (e.g., in ’t Zand & Weinberg 2010).

125

Alternatively, Weinberg & Bildsten (2007) attributethe ﬁrst brief dip as a pause between a spike due to a
shock breakout and a prematurely ignited type-I purehelium burst. However, Keek & Heger (2011) found thatthe heat released by the fallback of the photosphere after
the shock breakout may be suﬃcient to ignite the type-I
burst and that this occurs on time scales much shorter(10−5 s) than the dip time scale. After analyzing thespectral evolution during the dip, Keek (2012) concluded
that it could be attributed to photospheric expansion.
Keek (2012) also found that the precursor is more energetic than ordinary type-I bursts and, therefore, cannot
be powered solely by the burning of accreted helium.
They attribute the additional energy to the shock heating, supporting numerical models that predict that superbursts result in a detonation (Weinberg et al. 2006)
and shock that generates enough heat or overpressure to
power the precursor. A similar onset is observed in theother superburst detected with the PCA.
2.6. Burst quenchingSome superbursters are proliﬁc emitters of type-I bursts
at the time of the superburst, with recurrence times ofmerely a few hours. These include KS 1731-26 and 4U
1636-536. But it is noticeable that this emission of type-I
bursts is quenched for a considerable period of time after the superburst (Kuulkers et al. 2002; Cornelisse et al.
2002), namely about one month. Apparently, the superburst inﬂuences the nuclear burning for a considerably
longer time than when its emission is visible.
Measuring the duration of burst quenching is diﬃcult.
One needs a substantially high duty cycle to detect typeI bursts. This is not easily accessible with regular all-sky
monitors. Quench time measurements are very interesting, though. They provide a clean means to observe thetransition between stable and unstable nuclear burning(see Keek et al. 2012).
3. Theory
As with type-I bursts, the general picture is clear of
how superbursts come about (a deep thermonuclear shellﬂash), but there are two essential issues that need resolution.
3.1. Inferring basic physics parameters
Cumming & Macbeth (2004) and Cumming et al. 2006
showed that it is possible to infer from the light curve, inparticular the peak luminosity and the decay proﬁle, theignition column depth y12 = y/1012 g cm−2 and energy
yield E17 = E/1017 erg g−1 of the heated matter. The
higher the energy yield is, the higher the peak luminosity(up to the Eddington limit). The deeper the ignition is,
the longer the duration. An example of a ﬁtted superburst decay is shown in Fig. 4. Table 1 shows the valuesof these two parameters for many superbursts. Due toFig. 4. The decay of the superburst from 4U 1735-44 (see Fig. 1
ﬁtted with a model parametrized with y12 and E17 (see text andTable 1). From Cumming et al. (2006).

the lack of onset coverage, it is often diﬃcult to obtainreasonable constraints on y12 and E17 , particularly thelatter. Therefore, these numbers lack in many cases.
Keek et al. (2015) improved the light curve diagnostic
power by including the slope of the temperature relation with depth as a free parameter. This impacts inparticular the rise phase of the light curve.
The expected values for CNO burning and 3α process
are E17 = 64 and 15, respectively. This is 1 to 2 ordersof magnitude larger than the observed values for superbursts. The energy yield for nuclear burning of carbon
to iron-group elements is about 10. The superburst from
4U 1820-30 has the very same value.
3.2. What fuel is burning?
The longer duration points to a deeper ignition, a higher
density and a higher temperature than for type-I bursts.
This makes it easier to overcome the Coulomb barrierof heavier elements than hydrogen and helium. The
next most abundant element, certainly after the burningof hydrogen to helium through the CNO cycle and theburning of helium through the 3α process, seems carbon
and this is the fuel that was considered by Strohmayer& Brown (2002) and Cumming & Bildsten (2001).
The E17 values are more or less consistent with carbon burning. They range between 1.5 and 10. Thisimplies mass fractions of burnt carbon between 15 and100%. Part of the remaining composition are the heavy
isotopes needed for the right ignition depth (Cumming &Bildsten 2001) and whose photo-disintegration, furthermore, might contribute up to half the superburst energy,
possibly lowering the required carbon in the fuel (Schatz

126

et al. 2003).
3.3. Issue 1: how to obtain and maintain enough fuel?
If temperature becomes too high, the Coulomb barrierof carbon atoms is easier overcome by ambient protons
and alpha particles through the reactions 12 C(p,γ)13 N
and 12 C(α,γ)16 O, thus destroying the carbon. This caneasily happen during helium ﬂashes since the carbon andhelium are in the same layer. Thus, helium ﬂashes are
responsible for both the production and destruction ofcarbon. Generally, it is thought that the only mannerin which carbon destruction can be prevented is by preventing temperature to grow too large during 3α burning. This is only possible during stable helium burning. There is observational evidence for that, throughthe measurements of the burst α parameter (see §2.2.).
Also theoretically, improvements are made in the understanding. Stevens et al. (2014) proposes that the rpprocess consumes all protons before they get a chance tocapture on carbon. Keek & Heger (2016) introduces a
new regime of stable hydrogen burning that increases thetemperature in the H-depleted layer underneath somewhat, yielding higher 3α rates without going runaway.
3.4. Issue 2: how to ignite the fuel?
For diluted carbon mixtures, as suggested by the E17
values, it is diﬃcult to reach ignition conditions. The
temperature at a depth of y12 ≈ 1 is simply not highenough to overcome the Coulomb barrier. This temperature is primarily set by a heat ﬂow from the crust,
where pycnonuclear reactions and electron captures provide a heat source of Qb = 0.1 − 0.6 MeV/nucleon. The
power scales with the accretion rate. Cumming & Bildsten (2001) initially resolved this issue by proposing reduced conductivity of the ignition layer, thus reducing
the cooling of the layer. Then the crustal heating may
be enough to reach ignition temperatures, provided themass accretion rate is in excess of 0.1 times Eddington.
However, two recent ﬁndings aggravate the issue again.
First, the detection of superbursts from transient accretors, starting with 4U 1608-52 (Keek et al. 2008) andwith an extreme case of EXO 1745-248 (Serino et al.
2012; Altamirano et al. 2012). The accretion rate, averaged over the superburst recurrence time, is (much)
lower for transients so that crustal heating is accordingly
weaker. Second, recent calculations of neutrino coolingthrough the URCA process in the crust (Schatz et al.
2014) and the deep ocean (Deibel et al. 2016) show it tobe much more eﬃcient than previously thought.
The solution may come from ﬁnding shallow heatingprocesses, for instance due to rotational mixing and thus
deeper CNO burning (Keek et al. 2009) or freeze out ofheavy elements at the bottom of the ocean that induces
convection which heats up the superburst layer (Medin& Cumming 2011), but as yet the issue of unreachable
carbon ignition conditions remains.
4. Why study superbursts?
As may be clear from above, studying superbursts is interesting in its own right, but it is also very useful for
addressing a wide variety of scientiﬁc questions, for instance: 1) They are related to the same nuclear process(explosive carbon burning) that is thought to be responsible for type Ia supernovae. Understanding the ignition
will improve our understanding of type Ia SN ignition;2) Since the ignition is close to the neutron star crust,superburst characteristics depend on the thermal properties of the outer crust and, thus, provide a diagnosticof that crust (e.g., Cumming et al. 2006); 3) Superbursts
are sensitive probes of the neutron star spin and binary
orbit, through the detection of transient ms oscillations
during superbursts (Strohmayer & Markwardt 2002); 4)
Superbursts can be used as seismology probes of the neutron star interior, through the detection of oscillations
other than due to the spin (Strohmayer & Mahmoodifar2014); 5) Superbursts can be used as probes of the accretion disk. They irradiate and heat up the accretion
disk which can be observed through reﬂection features inthe superburst spectrum (e.g., Ballantyne & Strohmayer
2004; Keek et al. 2014). This irradiation happens bya simple spectral shape (black body) and over a broad
range of tractable temperatures; 6) The cooling of theneutron star envelope after a superburst has an observationally convenient time scale (days) to relatively easyprobe ignition conditions of hydrogen and helium burning (Keek et al. 2012).
5. Future
Only 11 superbursts have an observational coverage thatis better than 10%, strongly aﬀecting measurements ofpeak luminosities, quench times and onset proﬁles. Only
2 superbursts have been observed with an eﬀective areaof signiﬁcantly more than ∼100 cm2 , strongly aﬀecting
sensitive measurements of the spectrum (e.g., Keek et al.
2014) and interesting variability such as burst oscillations (Strohmayer & Markwardt 2002). Although MAXI
and other instruments are very useful in constraining
superburst recurrence times and ignition depths, it isobvious that much is to be gained from measurementswith higher duty cycles and larger sensitivity. This may
be obtained through 1) all-sky monitors that have large
2sky coverage and reasonable eﬀective area >∼ 100 cm in
the 1-10 keV band and 2) large X-ray telescopes withquick read-out times that spend substantial amountsof observing time on the population of potential superbursters (bursters with high α values) or can be quickly
brought on target after the onset of a superburst. One

127

concept platform where both of these types of instruments/observations are foreseen in a optimum manner is
LOFT (e.g., Feroci et al. 2016; in ’t Zand et al. 2015),
with spin oﬀ concepts eXTP (Zhang et al. 2016) and
Strobe-X (Wilson-Hodge et al. 2017). In the meantime, Astrosat (Agrawal 2006) and NICER (Gendreau
et al. 2012) are and will be valuable assets in the studyof superbursts if caught (see also Keek et al. 2016).
Acknowledgements. Laurens Keek and Motoko Serino
are thanked for useful comments on an early draft ofthis paper. This paper uses preliminary analysis resultsfrom the Multi- INstrument Burst ARchive (MINBAR),
which is supported under the Australian Academy of Sciences Scientiﬁc Visits to Europe program, and the Australian Research Councils Discovery Projects and Future
Fellowship funding schemes.
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129

MAXI Detections of Superbursts– Updates on Long X-ray Bursts –Motoko Serino1 and Wataru Iwakiri1
1

MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(MS): motoko@crab.riken.jp
Abstract

In seven years, MAXI observed twelve long X-ray bursts, which were observed during at least twoconsecutive scans. We divided them into two classes according to their e-folding decay time: the bursts
with longer e-folding decay time (> 1 hour) are superbursts, and the others are intermediate duration
bursts. MAXI detected superbursts from two bright persistent source (Ser X-1 and 4U 1705−44), a dim
persistent source (4U 0614+091), and four transient sources (Aql X-1, SAX J1747.0−2853, EXO 1745−248,
and SAX J1828.5−1037). Two of the superbursts from transient sources occurred at the time of very low
accretion rate, which may be before the outbursts. Four out of ﬁve intermediate duration bursts had low(< 2% of the Eddington limit) persistent ﬂux. Four intermediate duration bursts are from ultracompact
X-ray binary. These characteristics are common to the intermediate duration bursts which were observed
before the MAXI era. We found possible anti-correlation between the decay time and the peak ﬂux.
Superbursts have low peak ﬂuxes (10–50% of the Eddington limit) and intermediate duration bursts havethe peak ﬂuxes of about the Eddington limit.
Key words: stars: neutron — X-rays: bursts — MAXI

1. Introduction
MAXI scans about 85% of the whole sky in 92 minutes.
It is eﬃcient for searching for transient events with long(>92 min) duration such as long X-ray bursts. As ofthe end of 2016, we found twelve long X-ray bursts in
the MAXI data. Table 1 shows the list of long X-ray
bursts. Eight of them overlap with the list in in ’t Zandet al. (2017). Most of them are described in the previous
papers (Serino et al. 2016; Serino et al. 2017). Three
new events were found after Serino et al. (2016), whichwere IGR J17062−6143 (Negoro et al. 2015; Iwakiri etal. 2015; Keek et al. 2017), Ser X-1 (Iwakiri et al. 2016),
and 4U 1705−44 (Iwakiri et al. 2016). Details of Ser X1 and 4U 1705−44 are also introduced by Iwakiri et al.
(2017) in this proceedings. In this paper, we focus on
the persistent ﬂux of the long X-ray bursts and possible
anti-correlation between the decay time and the peakﬂux.
2. The persistent emissions of the long X-ray burst sources
MAXI monitors the persistent ﬂuxes of X-ray burst
sources. Figure 1 shows an image of galactic sourcesin 4–10 keV. The previously known superburst sources(dashed circles) and the sources in Table 1 (solid circles)

Table 1. A list of long X-ray bursts observed by MAXI (sorted by
date)

Name
4U 1820−30
SAX J1747.0−2853
EXO 1745−248
SAX J1828.5−1037
Ser X-1
SLX 1735−269
Aql X-1
4U 1850−086 (1)4U 0614+091
4U 1850−086 (2)IGR J17062−6143
4U 1705−44

date
2010-03-17
2011-02-13
2011-10-24
2011-11-12
2011-12-06
2012-12-06
2013-07-20
2014-03-10
2014-11-03
2015-05-09
2015-11-03
2016-10-22

decay time > 1h?
n
Y
Y
Y
Y
n
Y
n
Y
n
n
Y

10

130

Ser X-1
KS 1731-260

SLX 1735-269
EXO 1745-248
GX 3+1SAX J1747.0-2853

GX 17+2

40

30 4U 1850-086 20

Aql X-1

10

4U 1608-52

4U 1705-44

0

SAX J1828.5-1037

0

350

340

330

320

310

300

H 1636-536
4U 1254-690

4U 1735-444-10

4U 1820-30

IGR J17062-6143

Fig. 1. A MAXI image of galactic sources in 4–10 keV, based on data collection in 4.7 years. The previously known superburst sources(dashed) and the sources in Table 1 (solid) are marked with circles.

2
SAXJ1747

2-4 keV

0.3
0.2
0.1
0
-0.1
1
0.9
0.8
0.7
0.6
0.5
0.4
0.3
0.2
0.1
0

4-10 keV

-40

-20
0
20
time since superburst [day]40

Flux [photons cm-2 s-1]

0.4

Flux [photons cm-2 s-1]

-2 -1

Flux [photons cm s ]-2 -1

Flux [photons cm s ]0.5

AqlX-1

2-4 keV

1.5
1
0.5
0
-0.5
1.8
1.6
1.4
1.2
1
0.8
0.6
0.4
0.2
0
-0.2

4-10 keV

-40

-20
0
20
time since superburst [day]40

Fig. 2. Light curves of the outbursts of SAX J1747.0−2853 (left) and Aql X-1 (right). A point correspond to a scan. The short ﬂares attime= 0 are superbursts.

131

3. Peak ﬂux – decay time correlation
Figure 3 shows a scatter plot of the peak ﬂux and thee-folding decay time τ of long bursts. Since the peak
ﬂuxes are normalized with the observed maximum ﬂuxesof the normal bursts, which are an approximation of the
Eddington ratio2 . There are systematic uncertainties ofthe burst peak ﬂux, because of the limited scan time
window. Then we estimate the “possible maxima” by*1 http://maxi.riken.jp/top/*2 We do not have information of a normal X-ray burst fromIGR J17062−6143. We assume the observed peak ﬂux of thelong burst in our sample was the Eddington limit, because theinferred radius suggested a photospheric expansion (Keek et
al. 2017).

extrapolating the exponential function to the epoch of
the previous scan transit. The upper ends of the horizontal bars in ﬁgure 3 are these “possible maximum”.
There may be an anti-correlation between the decay
time and the peak ﬂux. Superbursts have low peak ﬂuxes(10–50% of the Eddington limit) and intermediate duration bursts have the peak ﬂuxes of about the Eddington
limit.
The authors thank Jean in ’t Zand for helpful comments on this paper and continuous collaboration withus.
10
4U0614

AqlX-1

EXO1745SAXJ1747

SerX-1
SAXJ1828

τ (hour)are marked. The former tend to be brighter than thelatter.
Two superbursts from SAX J1747.0−2853 and Aql X1 occurred during transient outbursts. The outburstsof these sources started ∼ 25 and ∼ 40 days before thesuperbursts. They are similar to 4U 1608−522, which
also occurred during the outburst and it started 57.6
days before the superburst (Keek et al. 2008).
No superburst was found from bright and persistentsources in the MAXI data before 2016. Serino et al.
(2016) mentioned that it should be a selection bias, because it is easier to ﬁnd superbursts from sources withweaker persistent emissions. Then we searched for superbursts from bright sources carefully, and found a superburst from Ser X-1. Another superburst, which wasobserved just before ”7 years of MAXI” conference, wasalso from a bright persistent source 4U 1705−44 (Iwakiri
et al. 2016; Iwakiri et al. 2017).
The MAXI light curve1 of SAX J1828.5−1037 stays
constant at ∼ 10 mCrab. We note that this emission
is dominated by the galactic ridge emission (see ﬁgure1). The light curves of the sources near the galactic center region (SLX 1735−269, SAX J1747.0−2853,
and EXO 1745−248) are also contaminated by nearby
sources. Therefore it is diﬃcult to estimate the persistent ﬂux from these sources.
We classiﬁed ﬁve bursts into the intermediate duration bursts, because they have exponential decay time
below 1 hour. Four out of these ﬁve had low (< 2% of
the Eddington limit) persistent ﬂux. Three sources (4U
1820−30, 4U 1850−086, and SLX 1735−269) with theintermediate duration bursts are (candidate) ultracompact X-ray binaries, while the compactness of the othersource (IGR J17062−6143) is not known. Ultracompact
nature and low persistent luminosity are common to thesources of intermediate duration bursts (Falanga et al.
2008). The persistent emission of 4U 1820−30 is exceptionally bright unlike other sources of intermediate
duration bursts.

4U1705

1

SLX1735
4U1850(2nd)

4U1820

IGRJ17062
4U1850(1st)

0.1
10

100

1000

Peak flux normalized with maximum flux of normal bursts (%)

Fig. 3. A scatter plot of peak ﬂux and e-folding time of long bursts.
Each peak ﬂux in the plot are normalized with the maximum
ﬂux of the normal bursts of the same source. Three new events
discovered after Serino et al. (2016) are labeled with larger font.
Note that the horizontal bars are not statistic errors. The lower
ends of the ﬂuxes are the observed ones, and the upper ends are
possible maximum ﬂuxes (see text).

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132

133

MAXI/GSC observations of new superbursts from Ser X-1
and 4U 1705−44
Wataru B. Iwakiri,1 Motoko Serino,1 and Jean in ’t Zand2

1

2

RIKEN MAXI team, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
SRON Netherlands Institute for Space Research, Sorbonnelaan 2, 3584 CA Utrecht, the Netherlands
E-mail(WI): wataru.iwakiri@riken.jp
Abstract

Two hours-long ﬂares were detected with MAXI/GSC, each from a type-I X-ray bursting low-mass
X-ray binary. The ﬁrst one was observed on December 12th, 2011, from Ser X-1. The second one wasobserved on 2016 Oct 22 from 4U 1705−44. On the basis of the results from the timing and spectral
analysis, we concluded that the ﬂares are the fourth superburst from Ser X-1 and ﬁrst superburst from 4U
1705−44, respectively.
Key words: thermonuclear burning — stars: neutron

2nd
3rd

5×10 42 − 4 keV
Photons/sec/cm2

0.6
0.4
0.2

4− 10 keV
Photons/sec/cm2

0
0.6
0.4
0.2
0
5.59×1045.5901×104
MJD

5.5902×104

5.5903×104

NEW

and simultaneously ﬁtted it with the burst spectrum extracted from each scan. The spectrum of each of the ﬁrst
three scan transits is well ﬁtted by an absorbed black+0.3
body model with temperatures kT = 2.1+0.3−0.2 , 1.5−0.2 and

0.2

1st

0.8

Fig. 2. The lightcurves from Ser X-1 observed by MAXI/GSC in the2 – 4 keV (upper panel) and 4 – 10 keV (lower panel).

0.1

Flux (Crab)

0.3

1. Superburst candidate from Ser X-1
Ser X-1 is a persistent source whose persistent brightness
is about 0.15 – 0.25 Crab (Figure 1). Three superbursts
were reported previously from this source (Cornelisse etal. 2002; Kuulkers 2009). The new event was found in
MAXI/GSC observations on 2011 Dec 12 (MJD 55901).
The event was detected at 08:10 UT and lasted for atleast 4 hours with a fast-rise and exponential-decay light
curve (Figure 2).

5.2×10 4

5.4×10 4
MJD

5.6×104

Fig. 1. Long-trend light curve of Ser X-1 obtained by RXTE/ASM(Black : 1.5 – 12.0 keV) and MAXI/GSC (Red : 2.0 – 10.0 keV) in
Crab units. The arrows denote the detection times of superbursts.

To obtain the information on the spectral variation
during a burst, we perform the spectral analysis. Assuming that the persistent emission did not change during
the burst, we extracted the spectrum from 10 consecutive scans just before the burst as the persistent emission

1.1+0.6−0.4 keV at 08:10, 09:42 and 11:15 UT, respectively.
The hydrogen column density was ﬁxed at 0.5×1022
cm−2 (Bhattacharyya & Strohmayer 2007). The results prove cooling along the decay. The bolometricﬂux declined from (6.7 ± 1.1) ×10−9 to (1.9 ± 0.7)×10−9
erg/s/cm2 over the course of 3 hours. From the obtained
bolometric ﬂuxes, the e-folding decay time is 2.3±0.7
hours. Since the interval of each MAXI scan is 92 min
and the observation time of each scan transit is about 40s, the peak ﬂux could be larger than the observed values.
Taking into account the interval of each MAXI scan, the

0.3
0.2
0.1
0

Flux (Crab)

0.4

134

5×104

5.2×104

5.4×104
MJD (day)

5.6×104

5.8×104

Fig. 3. Long-trend light curve of 4U 1705−44 obtained by RXTE/ASM (Black : 1.5 – 12.0 keV) and MAXI/GSC (Red : 2.0 – 10.0 keV) in
Crab units. The arrow denotes the detection time of the superburst.

1.2
2 − 4 keV
Photons/sec/cm2

2. Superburst candidate from 4U 1705−44
4U 1705−44 is persistent source and the long-trend
lightcurve (Figure 3) clearly shows the low and high intensity state. Because of the high-α value observed fromthis source (α is the ratio of the average persistent ﬂux
times the average wait times between two type-I burst
and the burst ﬂuence), it was predicted that 4U 1705−44
is a prospective superburster by in  t Zand et al. (2003).
However, no superburst was observed from this sourceyet.
The new candidate superburst from 4U 1705−44 wasfound in observations on 2016 Oct 22 (MJD 57683). Figure 4 shows the burst lightcurves. The observed new
event detected at 20:21 UT lasted for at least 2 hours
with a fast-rise and exponential-decay light curve. Theenergy-resolved lightcurves show that the X-ray emission
softens throughout the decay. We performed the spectral analysis following the same method described in theprevious section. The spectrum obtained during the ﬁrst
and second scan transits are well ﬁtted by an absorbed
blackbody model with temperatures of kT = 1.8 ± 0.1
and 1.8 ± 0.2 keV. The hydrogen column density was
ﬁxed at 1.9×1022 cm−2 (Piraino et al. 2007). This resultshows that the blackbody temperature did not decreasesigniﬁcantly but is typical of superbursts. The bolomet-

ric ﬂux declined from (1.4 ± 0.1) ×10−8 to (0.7 ± 0.1)
×10−8 erg/s/cm2 over the course of 1.5 hours. The efolding decay time is 2.2 ± 0.6 hours. From these results
we conclude that this is the ﬁrst superburst observed
from 4U 1705−44.
Interestingly, their type-I burst activity depends on
the intensity state; type-I bursts are more frequently occurred in the low intensity state (Langmeier et al. 1987).
On the other hand, we detected the superburst candidatein the high intensity state. It might be related to the prediction that the carbon fuel of the superburst could bedestroyed in type-I X-ray burst (Woosley et al. 2004).

4 − 10 keV
Photons/sec/cm 2

extrapolating peak ﬂux is (7 – 13) ×10−9 erg/s/cm2 . On
the basis of these results, we conclude that this is the 4thsuperburst event from Ser X-1.
We also calculated the ignition column depth y andenergy release per unit mass E using the cooling model
given by (Cumming & Macbeth 2004). Assuming the
distance is 6 kpc, the derived y and E are (1.7 – 2.5)
×1012 g/cm2 and (1.4 – 6.5) ×1017 ergs/g, respectively.
Comparing with the 1st superburst results (Cumming
et al. 2006), the ignition depth is signiﬁcantly deeper.
These results are summarized in Table 1 in in  t Zand(2017). The recurrence time between the 3rd and the 4thsuperburst is 1178 days although there is a possibility of
missing events during data gaps.

1
0.8
0.6
0.4
0.2
1.5

1

0.5

5.7682×10 4

5.7683×104

5.7684×104
MJD

5.7685×104

5.7686×104

Fig. 4. The lightcurves from 4U 1705−44 observed by MAXI/GSC
in the 2 – 4 keV (upper panel) and 4 – 10 keV (lower panel).

References
Cornelisse, R., et al. 2002, A&A, 382, 174
Kuulkers, E. 2009, The Astronomer’s Telegram, 2140
Bhattacharyya, S., & Strohmayer, T. E. 2007, ApJL,
664, L103
Cumming, A., & Macbeth, J. 2004, ApJL, 603, L37
Cumming, A., et al. 2006, ApJ, 646, 429
in  t Zand, J. J. M., 2017, 7 yeas of MAXI (theseproceedings)
Langmeier, A., et al. 1987, ApJ, 323, 288
in  t Zand, J. J. M., et al. & Cornelisse, R. 2003, A&A,
411, L487
Piraino, S., et al. 2007, A&A, 471, L17
Woosley, S. E., et al. 2004, ApJS, 151, 75

135

BinaryX-ray Pulsars

136

137

Monitoring X-ray Binary Pulsars with Fermi/GBM
P. A. Jenke1 and C. A. Wilson-Hodge2
1
2

University of Alabama in Huntsville, Huntsville, AL 35812, USA
NASA, Marshall Space Flight Center, Huntsville, AL 35812, USA
E-mail(PJ): peter.a.jenke@nasa.gov
Abstract

The Fermi Gamma ray Burst Monitor (GBM) is a unique instrument that oﬀers the largest instantaneous ﬁeld of view of any hard X-ray instrument currently in operation. This capability along withexcellent timing resolution, makes it very successful at detecting rare transient events as well as providing long integration times for pulsed signal extraction. Even though GBM has a relatively modest size,
we are able to observe a typical accreting pulsar for over 40,000 seconds each day allowing us to makeprecise measurements of the source frequency and pulsed ﬂux for sources with a spin frequency between
0.001 and 2 Hz. These frequency measurements due to GBM’s excellent timing capabilities have given us
the capability to determine/update orbital ephemerides for many sources as well as monitor rare torquereversals in persistent (semi-persistent) sources such as EXO 2030+375. Continuous Time Tagged Eventdata, available since November 2012, allows the GBM pulsar monitor to track the frequency of even higher
frequency sources and we plan to make these histories available this year.
Key words: pulsars,neutron stars

1. Introduction
Using the Gamma ray Burst Monitor (GBM), the GBM
Pulsar Monitoring program (GPM) is designed to monitor the frequency and pulsed ﬂux of accreting pulsars
and detect new outbursts from transient pulsars withspin periods between 0.5 to 1000 seconds and in the energy range between 8 and 50 keV. GBM is not an imaging
instrument but has the advantage of observing the entire
unocculted sky at all times with the exception of brieftransverses of the South Atlantic Anomaly. This fullsky coverage enables long term monitoring of accreting
pulsars, allowing precise measurements of spin frequencies, pulsed ﬂux, orbital parameters enabling studies of
the ﬂow of angular momentum from the donor star onto
the neutron star and constraints on the magnetic ﬁeld.
The long source exposures possible with GBM are crucial for precise timing measurements, where the error on
pulse frequency and pulse frequency rate estimates scaleas T−3/2 and T−5/2 , respectively, where T is the observation duration.
The pulsar monitor has a sister program, the GBM
Earth Occultation Monitor (GEOM), which uses the fall
and rise of the count rate of known sources as they are occulted by the Earth’s limb and reappear again (WilsonHodge et al. 2012). This program is currently monitoring over 200 sources and is able to provide ﬂux histories
for these sources in the energy range between 8 and 300

keV. A spin-oﬀ program, the X-ray Burst (XRB) monitor uses the pulsar monitor’s data cleaning process to
identify thermonuclear type 1 XRBs from low Ṁ accreting neutron stars and has identiﬁed 752 photosphericradius expansion (PRE) burst in its ﬁrst three years ofoperation starting in March 2010 (Jenke et al. 2016).
1.1. Gamma Ray Burst Monitor
GBM consists of 12 NaI detectors with a diameter of
12.7 cm and a thickness of 1.27 cm and two BGO detectors with a diameter and thickness of 12.7 cm. The NaI
detectors have an energy range from 8 keV to 1 MeV,
while the BGOs extend the energy range to 40 MeV.
There are three continuous publicly available data types:
CTIME data which has 8 energy channels (8–1000 keV
for NaIs) with 256 ms timing resolution and is primarily used for localization and long event searches, CSPEC
data which has 128 channels (8–1000 keV for NaIs) with
4,096 s timing resolution and is primarily used for spectroscopy, and CTTE data which has 128 channels (8–1000 keV for NaIs) with time tagged photon events with
2 μs precision and is used for spectroscopy, timing, andshort interval searches. The pulsar monitor primarily
uses CTIME data but is expanding its capability withCTTE data which increases sensitivity for higher frequency pulsars and improves pulsed spectral resolution
for bright sources.

138

2. GBM Pulsar Monitor
The pulsar monitor consists of two eﬀorts: (1) The daily
blind search aims at detecting new pulsed transients andoutbursts from known transients. (2) The dedicatedsearch is currently monitoring the pulsed ﬂux and frequency of 39 sources and has pulsed ﬂux and frequencies histories for 36 detected sources. GPM is providing,through its website1 , quick-look estimates of pulsed ﬂux
and frequency for use in multi-wavelength observations.
The website is updated twice a week and announcementsof new outbursts from transient sources are published in
Astronomer’s Telegrams to facilitate follow-up observations.
2.1. Pulsed Searches with GBM
The analysis of GBM observations of pulsars presents
two main challenges: the background rates are normallymuch larger than the source rates, and have large variations; and the responses of the detectors to a source arecontinuously changing because of Fermi’s ever changing
orientation. GPM meets these challenges with an empirical background subtraction and extensive on the ﬂy
calculations of detector response matrices. The initial
step in our analysis consists of screening the NaI detector count rate data to remove phosphorescence events,
gamma-ray burst and particle events and intervals of
rapid spacecraft rotation. Then an empirical backgroundmodel is ﬁt to the rates and subtracted oﬀ, after which
the rates are combined over detectors in a way that results in an estimate of the variable part of the source ﬂux.
From these ﬂux variation estimates we conduct searches
for pulsations, and monitor detected sources.
2.2. Empirical Background Subtraction
The screened rates in each channel of the 12 NaI detectors are ﬁt with a model with that includes components
for a small number of bright sources and a stiﬀ empirical model that contains the low frequency component of
the remaining rates. The source models are computedfrom a model ﬂux spectrum using time dependent response matrices. One ﬁt is made for each channel with
the normalization factors for the bright sources common
to all detectors, but with the stiﬀ model parameters (astatistically constrained spline) independent for each detector. The model accounts for the occultation steps forthe bright sources and the variation of the backgrounddue to spacecraft motion while upon subtraction retains
the higher frequency signals of the frequency modulated
sources.
2.3. Daily Blind Search
Using the background subtracted data, time dependent
responses and a generic pulsar spectrum (power law with
*1 https://gammaray.msfc.nasa.gov/gbm/science/pulsars.html

a high energy cut-oﬀ), source ﬂuxes (8–50 keV) are determined for 26 directions along the Galactic plane (15degree cadence) plus the LMC and SMC. The source
ﬂuxes are Fourier transformed and searched for signiﬁcant pulsations. The peak frequencies and directions are
determined by interpolating over the two dimensional
grid of Fourier power and directions. A catalog of accreting pulsars with the last known spin frequency is compared with the peak frequency and direction in order to
identify the pulsed signal. Signiﬁcant unidentiﬁed pulsed
signals are ﬂagged for additional analysis. Additional
analysis is limited by GBM’s large location uncertainty(a few degrees) but typically includes searching for signiﬁcant pulsations from a grid of directions that include
higher latitudes and searching for new ﬂaring pulsars in
the MAXI X-ray monitor (Matsuoka. et al. 2009) and
Swift BAT Hard X-ray Transient Monitor (Krimm et al.
2013).

2.4. Dedicated Search
In the dedicated search, source ﬂuxes (8–50 keV) aresimilarly determined for the speciﬁc monitored pulsar
position and energy spectrum of the source of interest.
Appropriate lengths of data (usually a minimum of 5–
10 times the pulse period) are ﬁt to a Fourier expansionin pulse phase using a course phase model. Times arebarycentered using the JPL Planetaria ephemeris DE200(Standish 1990) and, if an ephemeris is available, corrected for orbital motion. Fourier components for shortintervals (suitable for the source but usually a few days)
are combined and searched for signiﬁcant pulsations in
up to three harmonics over a small range of frequency
and possibly frequency derivative. This provides higher
sensitivity than the blind searches because of the smaller
search ranges and the use of source speciﬁc informationsuch as the location, orbital parameters and ﬂux spectrum.
GBM pulsar monitor currently has frequency histories and pulsed ﬂuxes for 36 accreting binaries including 28 transients. Figure 2 lists the currently detectedtransients, most of which are Be X-ray binaries, andthe times which GBM detects them. Of particular interest is the recent outburst of SMC X-3 in July 2016(see Figure 1), currently the only SMC source GBM detects. EXO 2030+375 is GBM’s most persistent transient source. Until recently, the GBM pulsar monitor
has detected an outburst almost every periastron passage. The Be X-ray binary’s spin-up ﬂattened and thesource faded from detection for 5 orbital periods in April
2016. When it reemerged on MJD 57632, the source had
undergone a rare torque reversal. Surprisingly, after two
more orbits, EXO 2030+375 appears to have resumed its
previous spin-up state (see Figure 3).

139

GRO J1744-28
4U 1901+03
4U 0115+634
V 0332+53
GRO J1750-27
SMC X-3
2S 1553-542
XTE J1859+083
RX J0520.5-6932
IGR J18179-1621
GS 0834-430
IGR J19294+1816
XTE J1946+274
2S 1417-624
KS 1947+300
Swift J0513.4-65
EXO 2030+375
Cep X-4
GRO J1008-57
2S 1845-024
A 0535+26
MXB 0656-072
RX J0440.9+4431
XTE J1858-034
GX 304-1
SAX J2103.5+4545
A 1118-616
MAXI J1409-619

SMC X-3Spin Frequency (mHz)

128.6
128.5
128.4
128.3
128.2
128.1

0.4
(keV cm-2 s-1)

12-50 keV Pulsed Flux

128.0

0.3
0.2
0.1

55000

55500

0.0
57660

57000

57500

57680

Fig. 1. The top panel is the frequency history of SMC X-3 while
the bottom panel is the 12–50 keV pulsed ﬂux. The frequencies
have been corrected for orbital motion using the ephemeris from
Townsend et al. (2017). Arrows indicate upper limits and the
diagonal bars through the frequency points represent the best
frequency derivative.

2.5. Orbital Ephemeris Determination
Typically, after correcting the pulse arrival times for
Earth’s motion, one can use the doppler boosted frequencies of the pulses to ﬁt a model for the binary system fora source of interest. Material accreted onto a neutronstar surface or collected in an accretion disk threaded by
the neutron star’s magnetic ﬁeld transfers angular momentum to the neutron star. Disentangling this intrinsic
spin-up from the orbital signature is challenging. A solution to this problem is to model the intrinsic spin-up
using a proxy for the system’s X-ray luminosity. The
luminosity is a function of mass accretion which is related to the torque imposed on the neutron star. Rappaport & Joss (1977) showed that, at high luminosity,
the intrinsic spin-up, ν̇, is proportional to L6/7 whenaccretion is mediated through a disk. The proportionality constant is a function of mass, radius, moment ofinertia, distance and magnetic ﬁeld of the neutron staralong with a few parameters describing accretion andemission eﬃciency. When connecting multiple outbursts
with the same torque model, it is necessary to include a
spin-down term to account for angular momentum losses
during quiescence (see Figure 4).
A search for frequency and frequency rate is performed
using pulse proﬁles from GBM folded over a multi–day
interval. Each frequency epoch is chosen as the mean
exposure–weighted observation time. The epochs arebarycentered using the JPL Planetary ephemeris DE200(Standish 1990). BAT survey data (15–50 keV) or MAXI
X-ray data (10–20 keV) for the individual source is used
as a proxy for the source luminosity and to model theFig. 2. The GBM detected transients are listed vertically while thetimes where they are detected are indicated by the blue bars.
EXO 2030+375
Spin Frequency (mHz)57640
Time (MJD)

24.22

24.21

24.20

24.19
0.5
(keV cm-2 s-1)

57620

12-50 keV Pulsed Flux

57600

56000
56500
Time (MJD)

0.4
0.3
0.2
0.1
0.0
55000

55500

56000
56500
Time (MJD)

57000

57500

Fig. 3. The top panel is the frequency history of EXO 2030+375 while
the bottom panel is the 12–50 keV pulsed ﬂux. Arrows indicate
upper limits. Notice the torque reversal and recovery around MJD
57300.

intrinsic spin-up rate. This spin-up model along with theline of sight delay associated with the binary orbit from
Deeter et al. (2003) is used to model the barycentered arrival times. Minimization of the χ2 ﬁt of the barycentricfrequencies and the count rates (Swift/BAT or MAXI)
is performed using the Levenberg-Marquart method andis of the form:χ2 =i (fi − (ν˙i (1 − βi ))+ i (Xi − ν˙i /m)22

; i = 0, n − 1(1); i = 0, n − 2(2)• βi is the orbital redshift factor at time ti which is afunction of the orbital elements.
• fi is the measured barycentric frequency at time ti
which is the frequency epoch of the search interval140   
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Fig. 4. The top panel is the frequency history of 4U 0115+634 while
the bottom panel is the 8–25 keV pulsed ﬂux. The frequencies
have been corrected for orbital motion. There are two giant outbursts in June 2011 and October 2015 where there is signiﬁcantspin-up do to accretion torques. After the giant outbursts, thepulsar spins down until a regular outburst occurs and there isagain spin-up do to accretion torques although the spin-up is not
as severe. Even between the normal outbursts there is signiﬁcant spin-down (red arrows) which is modeled as 1 or 2 degree
polynomial when ﬁtting the frequencies to an orbital model.

i. Each epoch, ti , is chosen as the mean exposureweighted observation time.
• Xi is the average value of R6/7 between ti and ti+1
and R is the MAXI count rate.
• The model parameter νi is the orbitally corrected
frequency at time ti and ν˙i is (νi+1 − νi )/(ti+1 − ti ).
The updated orbital model is used in a new search for frequencies and frequency rates recursively until the orbital
solution becomes stationary. Figure 5 demonstrates thistechnique. Notice how well the torque model follows thetransition between a normal and giant outburst around
MJD 56200.
2.5.1. Long Term Monitoring
Combining GBM data with data from the pulsar monitor’s predecessor which used data from the Burst andTransient Source Experiment (BATSE) on board theCompton Gamma Ray Observatory (CGRO), long term
pulsed histories spanning twenty years may be createdalthough there is a signiﬁcant time gap between CGRO
and the Fermi mission. Searches for orbital decay are
possible by determining the orbital epoch at intervals
over a long period of time. This is important for understanding how diﬀerent X-ray binary populations might

Fig. 5. Black points are the barycentered frequencies for GX 304−1.
The black curve is the best ﬁt torque model from the Swift/BAT
rates while the red curve is the orbital model combined with thetorque model. Notice how well the torque model represents thetorque during the onset of the giant outburst at MJD 56200.

be related through evolution. Orbital solutions are determined where the source of interest was bright for theavailable BATSE and GBM data. The Tπ/2 epoch is
expanded in orbit number, n,1
Tπ/2 = To + nPorb + n2 Porb Ṗorb ,2(3)

and the quadratic term determines the orbital decayof the source. The expected linear trend in the orbital epoch is subtracted and the residuals are ﬁt toa quadratic polynomial corresponding to the last termin Eq 3 (See Figure 6). Signiﬁcant orbital decay (Ṗ =(−9.74 ± 0.77) × 10−8 ) was detected in the wind fed highmass X-ray binary OAO 1657−414 (Jenke et al. 2012).
3. GBM Earth Occultation Monitor
The Earth occultation technique (EOT) measures thechange in ﬂux when a source enters (or exits) Earthoccultation (see Figure 7). The occultation time must
be known so a catalog of sources is used and updated
regularly. The changing GBM detector response, atmospheric transmission, the source spectrum, and othersources also occulting within a four minute ﬁt window(including the ﬂaring Sun) are taken into account. The
EOT is the only technique currently available that enables GBM to monitor persistent and transient galacticand extragalactic sources on timescales of days to years.
GEOM is currently monitoring 247 sources and light
curves (12–300 keV in four energy bands) and ﬁts ﬁles
(12–1000 keV in six energy bands) are available on

141

Fig. 6. Diamonds are orbital epochs for each ephemeris with a linear
trend removed. The curve is a quadratic ﬁt to the epochs. The
residual epochs to the left are calculated from BATSE data while
the residual epochs on the right are calculated from GBM data.

Fig. 8. GEOM results monitoring the ﬂux for Cygnus X-1 shown for
12–25, 25–50, 40–100, 100–300, 300–500 keV energy bands. The
vertical colored lines indicate times of state changes.

4. GBM XRB Monitor

Fig. 7. Example of an occultation ﬁt window. The occultation stepfor the source is clearly seen at t=0. The best model for the ﬁtwindow is shown by the solid curve.

the web site2 . GEOM has conﬁdently detected over
100 of its monitored sources with 9 sources being detected (> 5σ) above 100 keV (Case et al. 2011). The
sources detected above 100 keV include 7 persistent and
2 transient sources: the Crab, the BHs Cyg X-1, Swift
J1753.5−0127, GRS 1915+105, 1E 1740−29, and GRS
1758−258, the AGN Cen A, and the transient BHs XTE
J1752−223 and GX 339−4. We have also detected brief,bright events with single Earth occultations including
numerous solar ﬂares, and a brief episode of enhanced
persistent emission from SGR 1550−5418.
GEOM compliments the GPM by monitoring transient outbursts of Galactic accreting binaries and measure the total ﬂux of the source. GEOM can also monitor
spectral changes in these systems which may accompany
a change in the accretion state especially in bright blackhole binaries such as Cyg X-1 (see Figure 8).
*2 https://gammaray.msfc.nasa.gov/gbm/science/earth occ.html

Bright transient events are removed from the data prior
to ﬁtting the background in preparation for the pulsar
search. These events were cataloged for further analysis
under the GBM XRB Monitor. The events are localized by using the angular response of the NaI detectors
to reconstruct the most likely arrival direction based on
the diﬀerence in the background subtracted count rates
in the 12 NaI detectors which have diﬀerent sky orientations. The method is adapted from the method used for
GBM GRB localization (Connaughton et al. 2015). The
XRB localization uses the 12–50 keV rates and spectral
models appropriate for sources with softer energy spectra. Three cases are automatically identiﬁed: If the event
localization is consistent with being toward the Sun thenthe event is identiﬁed as being solar, If the orientationof the bright detectors are inconsistent with a single direction then the event is identiﬁed as a particle event.
Particle events tend to produce signal in all the detectors
or in detectors with opposite orientation. If the event localization is consistent with a single direction and is not
associated with the Sun then the event is considered acandidate XRB. These candidates undergo further analysis such as a reﬁned localization and spectral analysis
which occasionally resulted in a reclassiﬁcation and rejection of the candidate status reducing the ﬁnal number
of XRBs to 1084. Using spectral analysis, location andspatial distributions we classiﬁed the 1084 events into

142

References
Case, G. L. et al. 2011, ApJ, 729, 105
Connaughton, V. et al. 2015, ApJs, 216, 32
Deeter, J. E. et al. 2003, ApJ, 247, 1003
Jenke, P. A. et al. 2016, ApJ, 826, 228
Jenke, P. A. et al. 2012, ApJ, 759, 124Krimm, H. A. et al. 2013, ApJSS 209, 14
Linares, M. et al. 2012, ApJ, 760, 133
Matsuoka, M. et al. 2009, PASJ, 61, 999
Rappaport, S. and Joss, P. C. 1977, Nature, 266, 21
Standish, Jr., E. M. 1990, A&A, 233, 252
Townsend, L. J. et al. 2017, ArXiv e-prints, 1701.02336
Wilson-Hodge, C. A. et al. 2012, ApJS, 201, 33
Fig. 9. The red and green ﬁlled circles are the know type 1 XRBs.
The green ﬁlled circles are the XRBs that are suﬃciently isolated
from currently active sources that an association is possible. The
blue circles are the GBM detected events. The radius of the circle
is the approximate 1 sigma error region of the localization.

752 thermonuclear X-ray bursts (33 associated with 4U
0614+09), 267 transient events from accretion ﬂares and
X-ray pulses, and 65 untriggered gamma-ray bursts. See
Figure 9 for the distribution of detected type 1 XRBs.
All thermonuclear bursts have peak blackbody temperatures broadly consistent with photospheric radius expansion (PRE) bursts (Linares et al. 2012). We ﬁnd an
average rate of 1.4 PRE bursts per day, integrated overall Galactic bursters within about 10 kpc. These include
33 and 10 bursts from the ultra-compact X-ray binaries
4U 0614+09 and 2S 0918−549, respectively.

143

Luminosity and spin-up relation in X-ray binary pulsars
with MAXI/GSC and Fermi/GBM
Mutsumi Sugizaki1
1

RIKEN, MAXI Team, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(MS): sugizaki@riken.jp
Abstract

Relations between luminosity and pulse-period changes in X-ray binary pulsars are studied with lightcurve and pulse-period data obtained by the MAXI/GSC all-sky survey and the Fermi/GBM pulsar project,both covering over 7 years from 2009 to 2016. To examine the validities of accretion-torque models, weselected 12 Be binary pulsars whose stellar companions are certainly identiﬁed by optical observations
and surface magnetic ﬁelds are determined from the cyclotron resonance scattering feature. The obtained
luminosities, L, and spin-frequency derivatives, ν̇, during the large outbursts, clearly shows positive correlations expected from the models, ν̇ ∝ L6/7 , in all the 12 sources. The ratios of their relative variation
amplitudes, ν̇/L6/7 , agree within a factor of ∼ 3 with those expected from the disk-magnetosphere interaction developed by Ghosh & Lamb (1979) if the typical neutron-star mass of 1.4 M and radius of 10
km are assumed. The dispersion of the data-to-model ratios, ∼ 3, among the selected sample is explainedby the errors on the assumed emission-beam fractions and the source distances.
Key words: stars: neutron – X-rays: binaries

1.

Introduction

X-ray binary pulsars (XBPs) are systems consisting ofmagnetized neutron stars and mass-donating stellar companions. Since the neutron stars are strongly magnetized, the matter ﬂows from the companion are dominated by the magnetic pressure inside the Alfven radius,
and then funneled onto the magnetic poles along themagnetic ﬁeld lines. The accretion matter also transfers
its angular momentum at the Alfven radius. Therefore,
the pulsar spin-up rate and the mass accretion rate, i.e.
the X-ray luminosity, are thought to be closely correlated(e.g. Ghosh & Lamb 1979a, 1979b, hereafter GL79). The
relation is considered to reﬂect the manner of accretion
onto the neutron star, and also quantitatively relate to
the neutron-star parameters including mass, radius, andmagnetic ﬁeld. Thus, it is very important.
In this paper, we study the relation using the longterm light curves obtained by the MAXI/GSC allsky survey and the period changes obtained from thearchived results of Fermi/GBM pulsar project. These
data, taken by their continuous monitoring since 2009 forover 7 years, enable us to investigate their time variations
over the entire outburst activities of transient XBPs.

2. Observation data
2.1. MAXI GSC
Since the MAXI (Monitor of All-sky X-ray Image; Matsuoka et al. 2009) experiment onboard the International
Space Station started in 2009 August, the GSC (Gas Slit
Camera; Mihara et al. 2011), one of the two MAXI detectors, has been scanning almost the whole sky every
92-minute orbital cycle in the 2–30 keV band. To obtainthe long-term luminosity variation of Be XBPs covering the outbursts as well as the intermission/quiescence,
we use archived GSC light-curve data in 2–20 keV band,
which are processed with a standard procedure (Sugizaki
et al. 2011) by the MAXI team and archived at MAXI
web site1 .
2.2. Fermi GBM
The GBM (Gamma-ray Burst Monitor; Meegan et al.
2009) onboard the Fermi Gamma-Ray Space Telescope,is an all-sky instrument sensitive to X-rays and gammarays with energies between 8 keV and 40 MeV. The Fermi
GBM pulsar project (Finger et al. 2009; Camero-Arranz
et al. 2010) provides results of timing analysis of a number of positively detected X-ray pulsars, including their
pulsation periods and pulsed ﬂuxes via the web site2*1 http://maxi.riken.jp/*2 http://gammaray.nsstc.nasa.gov/gbm/science/pulsars/

144

since the in-orbit operation started in 2008 July. We
utilized the archived pulse period data.
3. Analysis and results
3.1. Relation between X-ray intensity and pulse-periodchange in various XBP subclasses
XBPs are classiﬁed into several groups according to thetype of companion (e.g. Reig 2011). Among ∼ 100
XBPs in the Galactic XBP catalogs (e.g. Liu et al. 2006),
the majority are high-mass X-ray binaries. They are further divided into two subgroups, supergiant (SG) XBPs
and Be XBPs. Low-mass (LM) XBPs are rare. Only afew (∼ 3) are known in our Galaxy. Figure 1 shows variations of 2–20 keV photon ﬂuxes F2−20 obtained from
the MAXI GSC light curves and spin frequencies νs of
Fermi GBM pulsar data for 7 years (from 2009 August to2016 March) in representative XBPs (in table 1). These
proﬁles are characterized by the the XBP subclasses, as
described in the below,
1. SG XBPs
Figure 1 shows the data of Vela X-1 and GX 301−2
as the representatives. In SG XBPs, accretion matter are fed by stellar wind. Thus, they usually appear as the persistent X-ray sources, and often exhibit time variations coupled with the orbital phases
The νs change is complicated, presumably due to theeﬀect of the stellar wind.
2. Be XBPs
Figure 1 includes 12 Be XBPs that have been selected by the certain identiﬁcations by both X-ray
and optical observations (section 3.2.). In this system, mass accretions start when the neutron stars
pass through the circumstellar disks around the Bestars. Therefore, the sources appear as transients,
synchronized with the orbital cycle. During theoutburst activities, spin up episodes are clearly observed. This is considered to be due to the accretion
torque.
3. LM XBPs
The mass accretion in LM XBPs are caused by
Roche-Lobe overﬂow. Thus, the sources show
persistent X-ray emission and steady spin-periodchanges. Figure 1 present the results of 4U 1626−67,
which has been used in Takagi et al. (2016).
3.2.

Selection of Be XBP sample to examine accretiontorque model
As seen in section 3.1. and ﬁgure 1, Be XBPs are best to
examine the accretion-torque models. We search for the
analysis targets that fulﬁll the following conditions: (i)optical companion is certainly identiﬁed as Be star, (ii)orbital elements are well determined, and (iii) signiﬁcantoutbursts were observed by both MAXI GSC and Fermi
GBM. Table 1 lists the selected 12 Be XBPs with theircharacteristic parameters, which include source name,pulse period Ps , orbital period Porb , eccentricity e, projected semi-major axis ax sin i, spectral type of the stellar companion, source distance D, surface magnetic ﬁeld
B12 estimated from the cyclotron resonance scatteringfeature (CRSF).
3.3. Spectral analysis for bolometric correction
The factor of conversion from the observed photon ﬂux tothe source luminosity depends on the the emission energy
spectrum. Hence, we analyzed the GSC energy spectra,
assuming that the energy spectrum of each source does
not change signiﬁcantly over the outburst active periods.
Thus, for each source, the spectrum was averaged over all
outburst periods, which is deﬁned as the period wherein
the GBM data are available.
Figure 2 shows the 2–30 keV spectra of the 12 sources,
obtained with the GSC. The background has been subtracted, but the instrumental responses are inclusive. We
ﬁtted them with a typical model for XBPs, consistingof a high-energy-cutoﬀ power-law (PLCUT) continuum,
and a Gaussian for iron-K line emission at 6.4 keV (e.g.
Makishima et al. 1999; Coburn et al. 2002). The PLCUT
is represented by the photon index Γ, the cutoﬀ energy
Ecut . the folded energy Efold , and the normalization Aas
E − Ecut,FPLCUT (E) = AE −Γ exp −
Efold

(1)

Because of the limited energy resolution, the centroid
and the width of the Gaussian were ﬁxed at 6.4 keV and
0.1 keV, respectively. For the interstellar absorption,
photoelectric absorption factor by a medium with Solarabundances was multiplied.
The PLCUT model was accepted, within 90 % conﬁdence limits of statistic uncertainty, except for X
0331+53. In ﬁgure 2, the best-ﬁt models folded with the
instrument response, are shown together with the data.a
Table 2 summarizes the best-ﬁt model parameters.
The residuals of X 0331+53 bear an absorption feature
at around 25 keV. This must be the fundamental CRSF
detected in past outbursts (e.g. Makishima et al. 1990;
Nakajima et al. 2010). We then multiplied a cyclotron
absorption (CYAB) model (Makishima et al. 1990) to
the PLCUT model, to ﬁnd the ﬁt acceptable. The bestﬁt CYAB parameters (table 2 footnote) are consistent
with those obtained in the past outbursts.
Using the best-ﬁt models, we calculated the factor of
conversion fbol from the 2–20 keV photon ﬂux to the
0.1–100 keV ﬂux (∼ the bolometric ﬂux) corrected forthe interstellar absorptions. They are presented in table

145

SG XBP
1.4
1.2
1
0.8
0.6
0.4
0.2

GX 301-2

1
0.8
0.6
0.4
0.2
0

GBM pulsar spin freq. [mHz]

GBM pulsar spin freq. [mHz]0
3.529
3.52853.528
3.5275
3 527

LM XBP

1.2

GSC photons cm2 s-1

Vela X-1

1.47
1.468
1.466
1.464
1.462
1.46
1.458
1.456
1.454

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

GBM pulsar spin freq. [mHz]

1.6

GSC photons cm2 s-1

GSC photons cm2 s-1

1.8

0.45
0.4
0.35
0.3
0.25
0.2
0.15
0.1
0.05
0
130.34130.32130.3
130.28130.26130.2455200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

4U 1626-67

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

MJD

0.6
0.4
0.2

2.5
2
1.5
1
0.5

276.69276.688276.686276.684276.682

228.52228.515228.51228.505

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

0.15
0.1
0.05
0
107.79107.78107.77107.76107.75
107.74107.73107.72
107.71
107 7

0.4
0.3
0.2
0.1
0
81.35

81.34681.34481.34281.34
81.338

GSC photons cm2 s-1

0.05

KS 1947+300

0.3
0.2
0.1

0.6
0.4
0.2

53.3
53.28
53.26
53.24
53.22
53.2
53.18
53.16

10.7
10.69
10.68
10.67
55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

63.46
63.44

0.3

EXO 2030+375

0.25
0.2
0.15
0.1
0.05

24.22524.22
24.21524.21
24.20524.2
24.19555200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

5

MJD
2.5

A 0535+262

4
3
2
1
0

10.71

63.5
63.48

MJD

GX 304-1

2
1.5
1
0.5
0

GBM pulsar spin freq. [mHz]

0
10.72

63.52

MJD

GSC photons cm2 s-1

0.8

0.1
0.05

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

GBM pulsar spin freq. [mHz]

GSC photons cm2 s-1
GBM pulsar spin freq. [mHz]

GRO J1008-57

0.2
0.15

0

53.32

MJD

1

0.3
0.25

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

01.2

XTE J1946+274
0.35

0.35

0.4

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

1.4

MJD

GBM pulsar spin freq. [mHz]

0
57.28
57.26
57.24
57.22
57.2
57.18
57.16
57.14
57.12
57.1
57.08

124.4

MJD

GBM pulsar spin freq. [mHz]

GSC photons cm2 s-1
GBM pulsar spin freq. [mHz]

0.1

124.5
124.4555200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

0.5

2S 1417-624

124.55081.348

MJD
0.2

124.6

0.4

0.5 GS 0834-430

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

0.15

0.02

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

GBM pulsar spin freq. [mHz]

GSC photons cm2 s-1

0.2

GBM pulsar spin freq. [mHz]

GSC photons cm2 s-1
GBM pulsar spin freq. [mHz]

0.3

0.04

MJD

0.4

0.25

RX J0520.5-6932

0.06

55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

H 1553-542
0.35

0.1
0.08

0

228.525

GSC photons cm2 s-1

276.692

GSC photons cm2 s-1

276.694

GBM pulsar spin freq. [mHz]0

GBM pulsar spin freq. [mHz]

GBM pulsar spin freq. [mHz]0

V 0332+53

GSC photons cm2 s-1

3

4U 0115+63

0.8

GSC photons cm2 s-1

1

GSC photons cm2 s-1

GSC photons cm2 s-1

Be XBP

9.685
9.68
9.675
9.67
9.665
9.66
55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

3.642
3.64
3.638
3.636
3.634
3.632
3.63
3.628
55200 55400 55600 55800 56000 56200 56400 56600 56800 57000 57200 57400

MJD

Fig. 1. 7 year (2009 August – 2016 March) variations of 2–20 keV photon ﬂuxes obtained from the MAXI GSC light curves and pulse
frequencies from the Fermi GBM pulsar data in representative XBPs. The vertical dotted lines represent the epochs of periastron passages.

146

Table 1. Characteritic parameters of XBPs used in the present analysis.

Type

Source name

SG

Vela X-1
GX 301−2
4U 0115+63
X 0331+53
RX J0520.5−6932
H 1553−542
GS 0834−430
XTE J1946+274
2S 1417−624
KS 1947+300
EXO 2030+375
GRO J1008−57
A 0535+262
GX 304−1
4U 1626-67

Be

LM

Ps
(s)
283.5
681.6
3.6
4.4
8.0
9.3
12.3
15.8
17.5
18.8
41.3
93.7
103.5
275.5
7.7

Porb
(d)
8.96
41.47
24.3
33.9
23.93
31.34
105.8
172.0
42.2
40.4
46.0
249.5
111.1
132.2
—

e
0.0898
0.462
0.34
0.37
0.03
0.04
0.12
0.33
0.45
0.02
0.41
0.68
0.47
0.52
—

aX sin i( lt-s )
114
368
140
78
108
201
202
471
188
137
246
439
267
498
—

Spec.Type
B0.5 Ib
B1-1.5 Ia
B0.2 Ve
O8.5 Ve
O8 Ve
B1-2 V
B0-2 III-Ve
B01 IVVe
B1 Ve
B0 Ve
B0e
B0e
O9.7 IIIe
B0.7 Ve—D( kpc )
1.9 ± 0.23
7.0 ± 0.3
6.0 ± 1.550 ± 220 ± 4+1
5−2
8.7 ± 1.2
11+1
−9
10.4 ± 0.9
6.5 ± 2.5
5.8 ± 0.5
2.1 ± 0.5
2.4 ± 0.5—B12(1012 G)
2.2
3.4
1.4
2.8
2.8
2.4
—
3.1
—
1.1
—
6.8
4.2
4.8
3.3

that the emission is isotropic (beam fraction fb = 1),
L = 4πD2 fb fbol F2−20 .

10−2

-2

Counts cm s-1 keV

-1

10−1

−3

10

10−4

2

4U0115V0332

2S1417
KS1947

RXJ0520.5

EXO2030

H1553

GROJ1008

GS0834
XTEJ1946

A0535
GX304

5

10

20

(2)

We also determined spin-fequency derivative ν̇s of each
6-d interval by linear ﬁt to the Fermi-GBM νs data.
Figure 3 shows the relations between ν̇s and L, calledν̇s -L diagram, for the 12 sources. These diagrams clearly
reveal the positive correlations, indicating that the pulsars indeed spin up by the accretion torque. Furthermore, in a fair fraction of the 12 objects, the correlation
is close to a direct proportionality (ν̇s ∝ L). Given this,let us apply the theoretical GL79 model to the observedν̇s -L relations.
In the GL79 model that describes mass accretion onto
magnetized neutron stars via Keplerian disk, the theoretical spin-frequency derivative, ν̇GL in units of Hz s−1 ,is expressed as a function of the source luminosity L37
in units of 1037 erg s−1 by

Energy ( keV )2, together with their 68% uncertainties caused by theﬁtting errors.
3.4.

Luminosity - spin-up relation in Be XBPs

We calculated source luminosities L from the 2–20 keV
photon ﬂux F2−20 in the MAXI-GSC light curves in 1-dtime bin, fbol (in section 3.3.), D (in table 1), assuming

−3/7

−1
ν̇GL = 1.9 × 10−12 μ30 n(ωs )R6 M1.4 I45
L37 (3)
2/7

Fig. 2. X-ray spectra of 12 Be XBPs taken by GSC during their
outburst-active periods. Solid lines represent the best-ﬁt PLCUT
models.

6/7

6/7

where μ30 , R6 , M1.4 , and I45 are the magnetic dipole
moment of the neutron star in units of 1030 G cm3 , theradius in 106 cm, the mass in 1.4 M , and the momentof inertia in 1045 g cm2 , respectively. The factor n(ωs ) isa dimensionless torque given as a function of the fastness
parameter ωs , which is the ratio of the angular velocityof the pulsar rotation to that of the Keplerian disk at
the inner radius r0 ,In ﬁgure 3, the black solid, dashed, and dotted curves
represent the model relations of 4U 0115+63, XTE
J1946+274, and GRO J1008−57, respectively, calculated
from equation (3), employing the canonical neutron-star

147

Table 2. Best-ﬁt PLCUT spectral parameters and spin-up correction factor η.

Source ID
4U 0115+63
X 0331+53§
RX J0520.5−6932
H 1553−542
GS 0834−430
XTE J1946+274
2S 1417−624
KS 1947+300
EXO 2030+375
GRO J1008−57
A 0535+262
GX 304−1

NH1022 cm−2
0.7+0.5
−0.5
< 0.2
< 0.9
1.9+1.8
−1.7
< 0.7
< 0.6
1.6+1.2−1.4
0.7+0.4−0.4
0.9+0.7−0.7
1.6+0.3−0.4
1.0+0.2−0.3
2.2+0.3−0.3

Γ
0.5+0.1−0.1
0.3+0.0−0.0
1.0+0.2−0.1< 0.4
0.8+0.1−0.3
0.8+0.1−0.2
0.6+0.2−0.6
1.2+0.1−0.1
0.6+0.2−0.3
1.0+0.1−0.1
0.9+0.1−0.1
1.0+0.1
−0.1

Ecut
(keV)
8.9+0.5−0.5
8.6+0.7−0.9
11.4+2.1−3.1
7.5+2.5
−1.7
< 4.2
< 2.7
7.5+2.3−3.9
8.3+1.7−1.3
4.1+0.6−0.5
7.5+0.6−0.8
8.2+0.9−1.3
7.2+0.5
−0.5

Efold
(keV)
5.7+0.6−0.6
19.7+18.5−6.3
6.0+6.5−3.5
9.5+7.5
−3.0
38+47
−14
22+11
−6
16.5+9.6
−5.7
27+18
−8
10.7+4.4−2.8
13.1+1.6
−1.5
24+3
−3
13.3+1.2−1.2

EWFeK †(eV)
< 112
134+41−43
170+163
−167
< 236
< 105
< 137
< 186
106+78
−80
99+55
−57
114+46
−52
53+34
−45
96+37
−36

χ2ν (ν)fbol ‡

η1.41(30)
1.09(27)0.72(30)0.62(30)
1.27(30)0.61(30)0.74(30)
1.42(30)
1.81(30)
1.01(30)0.91(30)
1.70(30)

1.44+0.08−0.07
4.04+2.84−1.23
1.32+0.35−0.18
1.66+0.43−0.28
3.21+0.63−0.52
2.48+0.41−0.34
2.55+0.48−0.40
2.07+0.28−0.23
1.75+0.25−0.20
1.84+0.08−0.08
2.63+0.09−0.08
1.88+0.06
−0.06

0.67
0.12
1.25
1.61
1.04
1.50
4.26
2.13
0.42
0.45
0.78
0.36

All errors represent 90% conﬁdence limits of statistical uncertainty.
†: Equivalent width of iron K (6.4 keV) line.
‡: Units in 10−8 erg counts−1 .
+2.5
+0.8
§
: CYAB model is applied. The best-ﬁt parameters are Ea = 23.4+1.0−0.8 keV, W = 6.2−2.3 keV, and D = 2.0−0.5 .

ν̇s = η ν̇GL ,10− 10GL79 model
4U0115
XTEJ1946
GROJ1008

ν (Hz s-1)

parameters, R6 = 1, M1.4 = 1, I45 = 1, and assumingthat the axis of the magnetic dipole is aligned to therotation axis (μ30 = B12 R63 /2). The model can generally explain the slope of the data distribution, but not
the absolute ν̇s values. This means that the coeﬃcientof proportionality between ν̇s and L6/7 , namely the ratio ν̇s /L6/7 , is not consistent between the data and themodel.
The discrepancy in the ν̇s -to-L6/7 coeﬃcient is primarily attributable to errors on the assumed parameters M ,R, I, in the model equation (3). Another origin may reside in the assumption of the GL79 model, that the gravity working on the accreting matter becomes counterbalanced by the pulsar magnetosphere at the radius ofr0 = 0.52rA .
To better compare the data and the accretion-torquemodel, we introduce a correction factor η to the GL79
model as10− 11
4U0115
V0332
RXJ0520.5
H1553
GS0834
XTEJ1946
2S1417
KS1947
EXO2030

10− 12(4)

and ﬁtted it to the data leaving η free. The best-ﬁt valuesof η, listed in table 2, distributed from 0.39 to 4.4, about
an order of magnitude, except X 0331+53 which required
an exceptionally small value of η = 0.12.
Figure 4 show a histogram of the 12 best-ﬁt η values,
where logarithmic bins are employed because the errors
on η are mostly proportional to η themselves. The average and the standard deviation of log η among the 11
sources, with X 0331+53 excluded, are log η = 0.004
and σ(log η) = 0.32, respectively. This means that the
average η is 100.004  1.0, and the 1-σ range is given by

GROJ1008
A0535
GX304

1

10
LX (1037 erg s-1)

102

Fig. 3. ν̇-L diagrams of the selected 12 Be XBPs. Black solid, dashed,
and dotted lines represent the GL97 models for 4U 0115+63, XTE
J1946+274, and GRO J1008−57, respectively,

148

Mihara, T., et al. 2011, PASJ, 63, 623
Nakajima, M., et al. 2010, ApJ, 710, 1755
Reig, P. 2011, Ap&SS, 332, 1
Sugizaki, M., et al. 2011, PASJ, 63, 635
Takagi, T., et al. 2016, PASJ, 68, S13

5

4

3

2

1

0

10−1

1

η

10

Fig. 4. Histogram of the correction factors to the GL79 models, η,
per logarithmic units. 12 arrows at the bottom point the values
for the 12 sources, and the red arrow at the top represent the log
average.

a factor of 100.32 = 2.1.
4. Summary and discussion
To examine the validity of the theoretical model of the interaction between the pulsar magnetosphere and the accretion disk in XBPs, we analyzed the X-ray light curves
and pulse-period variations of the 12 Be XBPs whose
distance, orbital elements, and surface magnetic ﬁeld are
well determined. The X-ray intensities was derived from
the MAXI GSC data, and the timing information was derived from the Fermi GBM, both for more than 6 years
since 2009.
In all the 12 objects, the expected ν̇s versus L6/7 proportionality was conﬁrmed. Except in X 0331+53, thecoeﬃcient of proportionality between ν̇s and L6/7 was
conﬁrmed to agree, within a factor of 3, with that predicted theoretically by GL79. The large discrepancy
found in X 0331+53 is likely to arise from its distanceoverestimation.
The average of the 11 ν̇s versus L6/7 coeﬃcients isalmost identical to the GL79 prediction. The scatter bya factor of ∼ 3 can be explained by uncertainties in the
involved parameters, in particular, those in D and fb .
References
Camero-Arranz, A., et al. 2010, ApJ, 708, 1500
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149

MAXI/GSC observation of X-ray outbursts from Be/X-ray binary pulsars
Motoki Nakajima,1 on behalf of the MAXI team1

School of Dentistry at Matsudo, Nihon University, 2-870-1, Sakae-cho Nishi, Matsudo-City, Chiba, Japan
E-mail(MN): nakajima.motoki@nihon-u.ac.jp
Abstract

MAXI/GSC has been monitoring the X-ray activities of Be/X-ray binary pulsars for 7 years since the
MAXI operation started. During this period, over the hundred outbursts were observed, and 50 Atels
were issued based on the MAXI/GSC observation data. So far, two types of the X-ray outbursts (normal
and giant) have been observed. The observations of the consecutive X-ray outbursts from A 0535+26 andGX 304-1 found the systematic shift of the outburst peaks. This systematic shift can be explained by theprecessed Be disk model. Comparing with the previous results, we ﬁnd that these outburst phase shift
rate might depend on the orbital parameters.
Key words: pulsars: individual(A0535+26,GX304-1,GRO J1008-57) — X-rays: binaries

1. Introduction
Be/X-ray binary pulsar (BeXBP; Reig 2011 and references therein) composed of an O or B type main sequencestar with Balmar emission lines (Be star) and a highly
magnetized (≥ 1012 gauss) neutron star orbiting an eccentric system is a representative subgroup of the X-ray
transients. This source occassionally or regularly exhibita X-ray outburst lasting for a few days to months, due
to the mass accretion from a circumstellar disk aroundBe star. The class of the X-ray outbursts are generally
divided into two types regarding the orbital phase andits luminosity. A normal (type-I) outburst takes placeat or near periastron and its peak luminosity distributes
in a narrow range of 1036−37 erg s−1 . This outburst is
generally repeated by the orbital period. Whereas a giant (type-II) outburst exhibits a large X-ray luminosity(≥ 1037 erg s−1 ) at its peak and their orbital phases
out of the periastron (Priedhorsky & Terrell 1983; Reig2011).
So far, several all-sky X-ray monitor instruments
have been observing the X-ray activities of BeXBPs.
From 2009 August, Monitor of All-sky X-ray Image (MAXI;Matsuoka et al. 2009) onboard Internatioal
Space Station had started its operation, and the monitor data are archived via web page. Thanks to the highsensitivity of GSC and continuous observations, a lot of
X-ray outbursts have been observed and their data aresystematically catalogued in the data base. In this paper, we introduce a brief observational results, about peculiar periodicity of X-ray outbursts from BeXBPs. We
will report the analysis results of the BeXBP outbursts
with the past/present all-sky monitor data, centering on

MAXI data.
2. MAXI/GSC observation of X-ray oubursts
So far, over 100 outbursts from BeXBPs have been observed by MAXI/GSC. Figure 1 shows the representative
light curves of BeXBPs. We have issued some reports to
Astronomer’s Telegram and MAXI Mailing List about
these detected events. In the cases of GX 304-1 andGRO J1008-57, the detection of the onset of the giant
outbursts were reoprted on 2010 August (Mihara et al.
2010) and 2012 November (Nakajima et al. 2012), andits alert information led to the discovery of cyclotronlines by the Suzaku ToO observations (Yamamoto et al.
2011; Yamamoto et al. 2014). In addition to the above
urgent followup observations, we also issued some reports
about the ordinary or extraordinay X-ray outburst behavior of BeXBPs. In the following, we focus on thepeculiar outburst periodicity.
3. Outburst Periodicity
Based on the MAXI/GSC data, we have reported somepeculiar outburst periodicity. Usually, the normal outburst recurs synchronized with its orbital period. However, the representative BeXBP, A 0535+26 and GX
304–1, exhibited the consecutive X-ray outbursts whichwere not repeated by the orbital period.
The ﬁrst report by the MAXI/GSC observation of the
X-ray outburst from A 0535+26 was issued on 2009 (Sugizaki et al. 2009). The orbital parameters, including theorbital period of 111.1±0.1 day, eccentricity (0.47±0.02)
and the epoch of the periastron passage (MJD=53613.0),

150

10

10
SMC Xí31

0.1
10
4U0115+63

1

GX304í1

1

0.1
10

0.1
10
V0332+53(X0331+53)

0.1
10
LS V+44 17

1
0.1
10

A0535+26

1
0.1
10

GS1843í02

1
Flux(photon/cm2/s)

1
Flux(photon/cm2/s)

H1145í619

1

0.1
10

0.1
10
XTE J1858+034

1
0.1
10

XTE J1946+274

1
0.1
10

4U0728í25

1
0.1
10

KS1947+300

1
0.1
10

GS0834í430

1
0.1
10

EXO2030+375

1
0.1
10

GROJ1008í57

1
0.1

Cep Xí4

1
0.1

55000

56000

57000
MJD

55000

56000

57000
MJD

Fig. 1. The 2-20 keV light curves of the representative BeXBPs observed by MAXI/GSC.

were determined by Finger et al. (2006). This source had
started a renewed activity from 2008, and the source
exhibited quasi-periodic X-ray outbursts. Interestingly,
these outburst recurred with the period of 115 days, notthe orbital period of 111 days (Nakajima et al. 2014).
Thus, the outburst peak phase shifted steadily withthe change rate of 3.0 × 10−4 phase d−1 . To interpretthe outburst phase shift phenomenon, precessed Be disk
model was proposed (Moritani et al. 2013 and references
therein). According to this model, the Be disk precessewith the period of 8.7 years. This period is comparable
to the V/R variation periods of isolated Be stars, ∼ 7 yr
(Okazaki 1991 and references therein).
Another source which exhibited the outburst phaseshift observed by MAXI/GSC is GX 304–1 (Nakajima
et al. 2016). This source was reactivated from 2008 after
the 30 years quiescent state, and has exhibited the seriesof normal/giant X-ray outbursts. The systematic outburst phase shifts were observed between 2014 Augustand 2016 January. Figure 2 shows the normal outbursts
light curves of GX 304–1 which were divided into theorbital period of 132.189 day intervals (Sugizaki et al.
2015). During this period, the outburst onset phases
shifted steadily from ∼ 0.90 orbital phase to 0.02. The
phase shift rates of the outburst onset is estimated to be(2.5±0.1)×10−4 phase d−1 . Thus, it is revealed that theoutburst phase shift rates of two sources are comparable.
In addition to the above analysis results obtainedby the MAXI/GSC data, we next introduce the othersources which showed the systematic outburst phaseshift. There are two BeXBPs, GS 0834–430 and EXO
2030+375, which showed the outburst phase shift in
past. In the case of GS 0834–430, the consecutive Xray outbursts had been observed from 1991 to 1993 by
CGRO (Wilson et al. 1997). In this period, the phaseshift rate was estimated to be 2.4 × 10−3 phase d−1 .
On the other hand, EXO 2030+375 exhibited not only
the steady shift of the outburst peak phase but also theorbital phase jump in 1996 (Wilson et al. 2002). The
previous analysis reported that the phase shift rate wasdetermined to be 1.8 × 10−4 phase d−1 . Besides, Laplace
et al. (2017) suggested that there is a possible long-term(∼ 10 or ∼ 20 years) periodicity of the giant outburst
due to Kozai-Lidov mechanism in the circumstellar disk
around Be star.
According to above results, we conﬁrmed that at least
4 BeXBPs exhibit the systematic phase shift of the outbursts. Therefore, the orbital phase shift can be considered as the common phenomenon among the BeXBPs.
In the followings, we attempt to compare the results between each source.
According to the previous observations, it is turned
out that the outburst phase shift rate of GX 304–1 ((2.4 ± 0.10) × 10−4 phase d−1 for the outburst onset) issame order of the results of EXO 2030+375 (Wilson etal. 2002 ; (1.8 ± 0.4) × 10−4 phase d−1 ) and A 0535+26(Nakajima et al. 2014 ; (3.03 ± 0.06) × 10−4 phase d−1 ).
Whereas the case of GS 0834–430, it is revealed that its
rate (2.4 × 10−3 phase d−1 ) is signiﬁcantly higher thanthe others.

151

1

MAXI/GSC(2í20keV)
Swift/BAT(15í50keV)

0.01
OB phase shift rate (phase dí1)

Next, we have searched for a correlation between theorbital phase shift rates and the other orbital/Be-starparameters. The most probable one which relates tothe orbital phase shift rates is the orbital eccentricity.
Figure 3 shows relation between the eccentricity and theshift rates. In the low eccentric system, the truncation
radius of the Be disk is much larger than that of thehigh eccentric binaries (Okazaki and Negueruela 2001).
Thus it might be considered that the tidal eﬀect aﬀects
the orbital phase shift rates of the outbursts. However,with the small number of the data, we can not conculdethat the eccentricity is a key parameter of the orbital
phase shift of the outbursts. Further discussion will bereported in the forthcoming paper.

GS 0834í430

10í3A 0535+26
GX 304í1
EXO 2030+375

10í4

0.2

0.4

Fig. 3. The relation between eccentricity and outburst phase shift
rate.

MJD:56814í56945

0.1

Matsuoka, M. et al. 2009, PASJ, 61, 999

0.01

Mihara, T. et al. 2010, Atel, 2779

10í31

MJD:56946í57077

Moritani, Y. et al. 2013, PASJ, 65, 83

0.1

Nakajima, M. et al. 2012, Atel, 4561

0.01

Nakajima, M. et al. 2014, PASJ, 66, 9Flux (Crab)

0.6

eccentricity

10í3
1

MJD:57078í57209

Nakajima, M. et al. 2016, Atel, 8592

0.1

Okazaki, A. 1991, PASJ, 43, 7

0.01

Okazaki, A. T. & Negueruela, I. 2001, A&A, 377, 161

10í31

MJD:57210í57341

Priedhorsky, W. C. & Terrell, J., 1983, ApJ, 273, 709

0.1

Reig, P. 2011, Astrophys. Space Science 332, 1

0.01

Sugizaki, M. et al. 2009, Atel, 2277

10í3
1

MJD:57342í57474

Sugizaki, M. et al. 2011, PASJ, 63, 635

0.1

Sugizaki, M. et al. 2015, PASJ, 67, 7311

0.01

Wilson, C. A. et al. 1997, ApJ, 479, 388

10í31

MJD:57475í57606

Wilson, C. A. et al. 2002, ApJ, 570, 287

0.1

Yamamoto, T. et al. 2011, PASJ, 63, 751

0.01

Yamamoto, T. et al. 2014, PASJ, 66, 5910

10í3
í0.4

í0.2

0

0.2

0.4

orbital phaseFig. 2. The light curves of GX 304–1 shown for individual 132.189-d
orbital cycles and those ﬂuxes in Crab units. Red marks represent the data obtained by MAXI/GSC, and blue ones are theSwﬁt/BAT results. The origin of the orbital phase (= 0) corresponds to the periastron passage.

References
Finger, M. H. et al. 2006, in BAAS, Vol. 38, BAAS, 359
Laplace, E. et al. 2017, A&A, 597, 124

152

153

Cyclotron absorption lines in the era of Suzaku and NuSTAR
Gaurava K. Jaisawal and Sachindra Naik

Astronomy and Astrophysics Division, Physical Research Laboratory,
Navrangapura, Ahmedabad - 380009, Gujarat, India
E-mail(GKJ): gaurava@prl.res.in
Abstract
Cyclotron resonance scattering features or cyclotron absorption lines are unique features observed in thehard X-ray spectra of accretion powered X-ray pulsars with magnetic ﬁeld of the order of 1012 G. Detectionof these features enables us for the direct estimation of strength of the magnetic ﬁeld close to the neutronstar surface. Corresponding to magnetic ﬁeld of ∼1012 G, the fundamental cyclotron lines are expected
in 10-100 keV energy range with harmonics expected at multiples of fundamental line energy. However,
we detected ﬁrst harmonic of cyclotron line at less than twice of the fundamental line energy (∼1.7 times
the fundamental line energy) in Be/X-ray binary pulsar Cep X-4. With the broadband spectral capabilityof Suzaku and N uST AR observatories, we have investigated cyclotron resonance scattering features inseveral X-ray pulsars to understand the shape of the lines, width, magnetic ﬁeld mapping, anharmonicity
in the line energies and luminosity-dependent properties of cyclotron lines. The results obtained fromthese works and new detection of cyclotron line in unknown/poorly studied sources are presented in thispaper.
Key words: stars: neutron – pulsars : individual (4U 1909+07, Cep X-4, SMC X-2) – accretion,
magnetic ﬁeld – X-rays: stars

1.

Introduction

Accretion powered X-ray pulsars are known to be highly
magnetized rotating neutron stars. These sources are
powered by the accretion of matter from the binary
companion either through Roche lobe overﬂow or capture of stellar wind from the massive companion. The
broadband spectrum in 0.1-100 keV range showed thepresence of several features such as ﬂuorescence emission lines, soft X-ray excess, broad absorption like features known as cyclotron resonance scattering features(CRSFs). CRSFs are generally detected in the hard Xray spectrum of pulsars with magnetic ﬁeld in the orderof 1012 G (Mészáros 1992). These line-like features areformed due to the resonant scattering of photons withelectrons near the neutron star surface. Depending on
the magnetic ﬁeld strength, the energy states of the electron are quantized harmonically in Landau levels with anenergy separation of Ecyc =11.6 B12 × (1 + z)−1 (keV),
where B12 is the magnetic ﬁeld in the unit of 1012 G
and z is the gravitational red-shift. Detection of CRSFs
in the pulsar spectrum, thus, is a powerful tool for the
direct estimation of magnetic ﬁeld of the neutron star.
After the launch of Suzaku and N uST AR, the number of cyclotron line sources has been rapidly increased

and helped us in understanding the distribution of magnetic ﬁeld lines around the poles of neutron stars. To
date, detection of cyclotron lines are conﬁrmed in 30
sources. However, there are about 8 sources in whichthese lines are tentatively detected (Table 1). Apartfrom the fundamental line, harmonics of cyclotron lines
are also seen in some cases. Many interesting aspects related to the line shape, anharmonicity in coupling factor,
positive/negative dependence of line energy with luminosity have been studied in recent years. Among these,
we present some recent results on cyclotron line obtained
from the studies of X-ray pulsars such as 4U 1909+07,
Cep X-4 and SMC X-2 by using Suzaku and N uST AR
observations in this article.
2. Observations and Analysis
The broadband spectral studies of 4U 1909+07, Cep X4 and SMC X-2 were performed by using data from
Suzaku (Mitsuda et al. 2007) and N uST AR (Harrison
et al. 2013) observatories. 4U 1909+07 and Cep X4 were observed with Suzaku during 2010 Novemberand 2014 July, respectively. Observations of SMC X-2
were carried out with N uST AR at three epochs during 2015 October X-ray outburst. The energy spectra

154

CRSFs
(keV)

Ref.

1
2
3
4

Swift J1626.6-5156
XMMU J05∗
KS 1947+300
4U 0115+634

5
6
7
8
9
10
11
12
13
14
15
16
17
18
19
20
21
22
23
24
25
26
27
28
29
30

IGR J17544-2619
4U 1907+09
4U 1538-52
IGR J18179-1621
IGR J18027-2016
2S 1553-542
Vela X-1
V 0332+53
SMC X-2
Cep X-4
4U 0352+309
IGR J16393-4643
Cen X-3
IGR J16493-4348
RX J0520.5-6932
LS V+44 17
MXB 0656-072
XTE J1946+274
4U 1626-67
GX 301-2
Her X-1
MAXI J1409-619
1A 0535+262
GX 304-1
1A 1118-615
GRO J1008-57

10
10
12.5
14, 24, 36
48, 62
17
19, 4022, 4722
23
23.5
25, 5027, 51, 74
27
28, 45
29
29.3
30
30
31.5
32
33
36
37
37
39, 7344, 73, 128
45, 100
54
55, 112?
76

DeCesar et al. (2013)
Manousakis et al. (2009)Fürst et al. (2014)
Wheaton et al. (1979)
Ferrigno et al. (2011)
Bhalerao et al. (2015)
Rivers et al. (2010)
Rodes-Roca et al. (2009)
Li et al. (2012)
Lutovinov et al. (2017)
Tsygankov et al. (2016)
La Barbera et al. (2003)
Nakajima et al. (2010)
Jaisawal & Naik (2016b)
Jaisawal & Naik (2015b)
Coburn et al. (2001)
Bodaghee et al. (2016)
Burderi et al. (2000)
D’Aı̀ et al. (2011)
Tendulkar et al. (2014)
Tsygankov et al. (2012)
McBride et al. (2006)
Heindl et al. (2001)
Orlandini et al. (1998)
Suchy et al. (2012)
Enoto et al. (2008)
Orlandini et al. (2012)
Terada et al. (2006)
Yamamoto et al. (2011)
Doroshenko et al. (2010)
Yamamoto et al. (2014)

1
2
3
4
5
6
7
8

EXO 2030+375
GS 1843+009
2S 0114+650
GX 1+4
OAO 1657-415
4U 1700-37
4U 1909+07
LMC X-4

11? 61?
20?
22?, 44?
34?
36?
37?
44?
100?

Wilson et al. (2008)
Mihara et al. (1995)
Bonning et al. (2005)
Ferrigno et al. (2007)
Orlandini et al. (1999)
Jaisawal & Naik (2015a)
Jaisawal et al. (2013)
La Barbera et al. (2001)

XMMU J05∗ stands for XMMU J054134.7-682550

of pulsars were described with diﬀerent phenomenological models such as high energy cutoﬀ power-law model(highecut), cutoﬀ power-law model, Fermi-Dirac cutoﬀ
model (fdcut), negative and positive exponential cutoﬀ
power-law (NPEX) and more physical Comptonization
model such as CompTT. The spectral ﬁtting was carriedout using XSPEC package. In addition to continuum,
additional components such as partial covering, Gaussian functions for emission lines from ﬂuorescence emissions, CYCLABS or Gaussian absorption component for
cyclotron features were also used whenever required.

χ

Sourceχ Normalized Counts sí1 keVí1

1

Table 1. Up to date list of cyclotron line sources.

0.1
0.01
10í3
4
2
0
í2
í4
4
2
0
í2
í4
1

10
Energy (keV)

Fig. 1. The 1-70 keV energy spectrum of 4U 1909+07 ﬁtted withpartial covering NPEX model along with 6.4 keV iron line anda cyclotron line at 44 keV. Middle and bottom panels show thecontributions of the residuals to χ2 for each energy bin without
and with cyclotron line, respectively.

3. Results
3.1. Cyclotron line in 4U 1909+07
4U 1909+07 is a slow pulsar with spin period of 604 s(Levine et al. 2004). The optical companion of the Xray pulsar was found to be an OB supergiant star (Morel& Grosdidier, 2005). The orbital period of the system
was reported to be 4.4 d. The broad-band X-ray continuum of the pulsar was poorly studied in the past due
to the dearth of high quality data. We have studied
the 1-70 keV wide band spectrum of pulsar by using a
Suzaku observation. The partial covering NPEX andhigh energy cutoﬀ models were found to describe thecontinuum well. Apart from the continuum model, additional components such as a blackbody component forsoft X-ray excess and a Gaussian function for 6.4 keViron emission line were also required in the spectral ﬁtting. While ﬁtting the pulsar spectrum with standard
continuum models, an absorption-like feature was seenat ∼44 keV (see Fig. 1). We have investigated the spectrum with combination of several models to check thepossibility of model dependency of the feature. However, a broad absorption feature was clearly seen in thepulsar continuum in a model independent manner. We
identiﬁed this feature as the cyclotron absorption lineof the pulsar. We attempted to normalize the pulsar
spectrum with that of Crab pulsar which has a featureless continuum with photon index of 2.1. The crab ratio
clearly showed the presence of absorption feature in 40
to 50 keV range (see Figure-7 from Jaisawal et al. 2013).
This also conﬁrmed the presence of cyclotron feature in
4U 1909+07. The statistical tests such as F-test andrun-test were performed to investigate the signiﬁcance of
the absorption feature. Based on these analysis, we re-155

Normalized Counts sí1 keVí1

0.01

10í3

10í4

χ

Fundamental and ﬁrst harmonics of cyclotron line in
Cep X-4
Be/X-ray binary pulsar Cep X-4 was observed with
Suzaku during its 2014 June-July Type I X-ray outburst. We have used data from a ∼60 ks observation of
the pulsar in the declining phase of the outburst. Cep X4 is a pulsar which have a pulsation period of 66.33 s.
A fundamental cyclotron line was discovered in the pulsar at ∼30 keV with Ginga observation (Mihara et al.
1991). In this work, we report the detection of the ﬁrst
harmonic of the cyclotron line. The 1-70 keV spectrum
obtained from Suzaku observation was well describedwith standard continuum models such as NPEX, hecut,
CompTT along with partial covering component andGaussian functions for two emission lines at 6.4 and 6.9
keV. While ﬁtting, a strong absorption feature was detected at 28 keV which has been known as the cyclotronline in the pulsar. In addition to this, an absorption likefeature at ∼45 keV was also detected (see Fig. 2). Thisfeature was seen in the pulsar spectrum while ﬁtting withall the above continuum models. We have identiﬁed thisfeature as the ﬁrst harmonic of the cyclotron line though
it was detected at 1.7 times of fundamental line energy.
The statistical signiﬁcance of the absorption line was also
tested by using XSPEC script ‘simftest’. We found thedetection of feature with statistical signiﬁcance of >4σ.
Generally, harmonics of the cyclotron line is expected
at the multiples of fundamental energy. However, in caseof Cep X-4, we have detected the ﬁrst harmonic below
the ideal coupling factor 2. Anharmonicity in line energy
has also been seen in other pulsars, where line energies
ratio was found at above as well as below 2 (see Table 1 and references therein). Relativistic approximationof photon-electron scattering can contribute to the anharmonicity in cyclotron line energies (Mészáros 1992).
However, such low value of line ratio, as seen in Cep X-4,
is diﬃcult to explain in this approximation. There areother possibilities by which anharmonicity in cyclotronline energies can be understood. It is probable that thefundamental and harmonic lines are formed at two diﬀerent scale heights which can have diﬀerent optical depth.
As a result, anharmonicity in the line energy ratio canbe observed. Studies on cyclotron lines from Nishimura(2005) showed that increased in magnetic ﬁeld in line
forming region can produce the anharmonic cyclotronline energies. The eﬀect of viewing column/line forming
region at large angles can also be the cause (Nishimura2013). Alternatively, anharmonicity in line energies may
indicate the distortion or changes in the magnetic ﬁeld
geometry. However, we have not detected any signiﬁcant

0.1

5
0
í5

χ

3.2.

1

2
0
í2

χ

ported a probable detection of cyclotron line at ∼44 keV
in 4U 1909+07. Corresponding magnetic ﬁeld strengthof the neutron star was estimated to be 3.8×1012 G (Jaisawal et al. 2013).

2
0
í2

NPEX

NPEX*CYCLABSNPEX*CYCLABS*CYCLABS

1

10Energy (keV)

Fig. 2. The 1-70 keV energy spectrum of Cep X-4 ﬁtted with partial
covering NPEX model along with 6.4 and 6.9 keV iron lines andcyclotron lines. The second, third and fourth panels show thecontributions of the residuals to χ2 when pulsar continuum was
ﬁtted with a partial covering NPEX model without, with one andtwo cyclotron line components, respectively.

variation in cyclotron line parameters with the pulsarphases in phase-resolved spectroscopy (see Figure-4 from
Jaisawal & Naik, 2015b).
3.3. Discovery of cyclotron line in SMC X-2
SMC X-2 is a high mass X-ray binary pulsar with a spin
period of 2.37 s (Corbet et al. 2001). The orbital period of the binary system was estimated to be 18.4 d(Townsend et al. 2011). This source went in to an intense outburst in 2015 September and observed with major X-ray observatories. We have investigated spectral
properties of the pulsar with N uST AR and Swif t/XRT
at three diﬀerent luminosity epochs (>1038 erg/s). The
1-79 keV spectra were well ﬁtted with the traditionalpulsar continuum models such as NPEX, fdcut, cutoﬀ
power-law along with iron emission line. While ﬁtting,an additional absorption like feature was also detectedin the spectra at ∼27 keV for the ﬁrst time (see Fig. 3).
This feature was seen in a model independent mannerand identiﬁed as a cyclotron absorption line in the pulsar spectrum. Corresponding to the line energy, themagnetic ﬁeld of the neutron star was estimated to be
2.3×1012 G (Jaisawal & Naik, 2016b).
N uST AR observations also provided opportunity toprobe the luminosity dependence of the cyclotron linein this pulsar. As the observations were performed atthree diﬀerent luminosity levels, we detected a marginalchanges in the cyclotron line energy that was anticorrelated with the luminosity (Table 2 in Jaisawal &Naik, 2016b). These luminosity dependent variation in

156

cyclotron line in SMC X-2 along with a negative correlation between cyclotron line energy and luminosity anddiscussed these ﬁndings in term of changes in the lineforming region.

keV (Photons cmí2 sí1 keVí1)

0.1

0.01

10í3

10í4

χ

5
0

NPEX
NPEX*GABS

χ

5
0
1

2

5
10
Energy (keV)

20

50

Fig. 3. The 1-79 keV energy spectrum of SMC X-2 obtained from ﬁrst
N uST AR and Swif t/XRT at the peak of the outburst, ﬁtted
with partial covering NPEX model along with 6.4 keV iron line
and cyclotron line at 27 keV. The second and third panels show
the contributions of the residuals to χ2 when pulsar continuum
was ﬁtted with a partial covering NPEX model without and withcyclotron line component, respectively.

cyclotron line energy can be attributed to the changes
in the line forming region. There are a couple of sources
which show a positive correlation of cyclotron line energy with luminosity (e.g. GX 304-1; see discussion section in Jaisawal et al. 2016a). It is believed that these
correlations are attributed to a critical luminosity. The
pulsars below the critical luminosity show a positive correlation and vice-versa. In case of SMC X-2, the source
luminosity (>1038 erg/s) was in the super-critical regime
where the radiation pressure dominates and a shock can
be formed above the neutron star surface (Becker et al.
2012). As the luminosity increases, the line-forming region is expected to shift in the accretion column. This
results a negative correlation between cyclotron line energy and luminosity.
4. Summary
We have reported the detection of cyclotron line(s) inpulsars such as 4U 1909+07, Cep X-4 and SMC X-2
by using high quality data from Suzaku and N uST AR
observatories. We have also probed various phenomena
such as anharmonicity in line energies, luminosity dependent characteristics as well as mapped the magnetic
ﬁeld around pulsars. The presence of anharmonic lines
in Cep X-4 may indicate that the fundamental and ﬁrst
harmonic of cyclotron line were not forming at the same
scale. It is also possible that the line-forming regions are
viewed at larger angles. For the ﬁrst time, we detected aReferences
Becker, P.A. et al. 2012, A&A, 544, A123
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157

X-ray emission of the Be star/pulsar system PSR B1259−63
at non-periastron phases
Yuki ONO,1 Yoichi YATSU,1 Mikio MORII,2
Nobuyuki KAWAI1 and the MAXI team1

Department of Physics, Tokyo Institute of Technology, 2-12-1, Ohokayama, Meguro, Tokyo, Japan
2
The Institute of Statistical Mathematics, 10-3-D414, Midoricho, Tachikawa, Tokyo, Japan
E-mail(YO): onoy@hp.phys.titech.ac.jp
Abstract

We report on the MAXI observation of the ﬁeld of PSR B1259−63. X-ray ﬂares have been observedat orbital phases far from the periastron. However, IGR J13020−6359 is located at 0.16 deg from PSR
B1259−63, making it diﬃcult to identify the sources of the ﬂares. In this study, we ﬁt the X-ray images
with the Point Spread Function of MAXI GSC to determine the position of the X-ray ﬂares accurately.
As a result, we can distinguish the ﬂares of PSR B1259−63 from those of IGR J13020−6359; three ﬂares
out of ﬁve ﬂares are likely to be from PSR B1259−63. We suggest that its Be star companion temporarily
releases a large amount of gas in directions other than the circumstellar disk, resulting in increases of thestellar wind density around the pulsar at phases far from the periastron.
Key words: pulsars: individual(PSR B1259−63) — X-rays: binaries

1. Introduction
PSR B1259−63/SS 2883 is a binary system in which apulsar orbits around a Be star with a long orbital period
of about 3.4 year. It is argued that the system is nonaccreting and a shock is formed between the pulsar wind
and the Be star wind, resulting in particle accelerationand non-thermal high energy emission near periastronup to TeV gamma-ray bands (Tavani et al. 1997).
2. Observation
Figure 1 is the X-ray light curves of the ﬁeld of PSR
B1259−63. According to Figure 1, X-ray ﬂares havebeen observed at orbital phases far from the periastron.
However, the resolution of MAXI GSC is typically on
the order of 1.5 deg, while a nearby transient source IGR
J13020−6359 is located at 0.16 deg from PSR B1259−63,
making it diﬃcult to identify the sources of the ﬂares.
The source of Flare 1 is identiﬁed as IGR J13020−6359
by Swift/XRT observation (Kong et al. 2010) and thesource of Flare 2−5 has not been identiﬁed due to lackof concurrent XRT observation.
3. Analysis and results
In order to identify the source of the ﬂares, we determined the position of the X-ray ﬂares with using the PSF
ﬁt tools (Morii et al. 2016). We formed a PSF modelFig. 1. MAXI/GSC and Swift/BAT light curves of the ﬁeld of PSR
B1259−63. The upper panel shows the MAXI GSC light curve
and the lower panel shows the Swift BAT light curve. Periastrons
and ﬂares are highlighted with shades.

consisting of a ﬂare source, catalogued X-ray sources
excluding PSR B1259−63 and IGR J13020−6359, andbackground photons on a detector coordinate of MAXI
GSC. We then allowed the position of the ﬂare source
to move and determined its position to maximize thelikelihood function. Figure 2 is the results of the PSF
ﬁt determinating the position of the X-ray ﬂares. Fig-

158

Flare ID
Flare 2
Flare 3
Flare 4
Flare 5orbital phase
φ
0.085
0.641
0.962
0.148

Flux (1-20 keV)
×10−10 erg cm−2 s−1
2.81+0.51−0.60
4.76+0.50−0.71
7.41+0.50−2.31
3.83+0.21−0.73

L1−20keV×1035 [erg cm−2 s−1 ]1.77+0.32−0.37
3.00+0.32−0.45
4.67+0.32−1.46
2.42+0.13−0.46

Table 1. A list of the ﬂux and the luminosity of the observed ﬂares. We assumed that PSR B1259−63 is located at a distance of 2.3 kpc.

Fig. 3. The calculation results of the X-ray luminosity.

ξ
0.03
0.05
0.08
0.1
Fig. 2. The calculated position of the ﬂares with using the PSF ﬁttool. The cross marks show tha calculated position. The dashedlines show that the statistical 90 ％ conﬁdence level. The solidlines show that the 90 ％ conﬁdence level taking the systematic
error of MAXI GSC into account.

ure 2 shows that the source of three ﬂares out of four
ﬂares are consistent with PSR B1259−63 (Flare 2, 3, 5).
The position of the ﬂare 4 is near IGR J13020−6359 but
contains the coordinates of both IGR J13020−6359 and
PSR B1259−63. Therefore, we included the ﬂare 4 indiscussion.
4. Discussion
We assumed that synchrotron cooling dominates theemission of these ﬂares from the shocked wind and thatthe shocked electron pairs have a power-law distribution
N (γ) = Kγ −p (γmin < γ < γmax ). We then calculated
the X-ray luminosity emitted from the shock-accelerated
electron/positron (Murata et al. 2003). The parameters
we used for the calculations are the minimum Lorentz
factor γmin , the power law index p, the conversion eﬃciency from the pulsar spin-down luminosity to the relativistic electron/positron pairs a , and the ratio of theshock distance from the pulsar to the binary separation

Flare 2
46.7
16.1
5.92
3.63

Flare 3
161.5
55.3
20.4
12.1

Flare 4
12.0
4.13
1.51
0.929

Flare 5
108.3
37.6
13.6
8.23

Table 2. The mass loss rate of the sphericlly symmetric Be staroutﬂow Ṁ with each ξ. The mass loss rates Ṁ in the table isgiven in units of 10−6 M /yr.

ξ. We found that the observed ﬂux can be explained
when a is ∼0.8 and ξ is <0.1 (Figure 3).
We suggest that the Be star temporarily releases alarge amount of gas in directions other than the circumstellar disk, resulting in increases of the stellar wind density around the pulsar at phases far from the periastron.
When we assume that the Be star wind is sphericallysymmetric, such a small shock distance is possible whena mass loss rate of the Be star Ṁ is 10−4 − 10−6 M /yr(Table 2).
References
Sugizaki et al. 2011, PASJ, 63, 635
Mihara et al. 2011, PASJ, 63, S623-S634
Tavani & Arons 1997, ApJ, 477, 439-464
Murata et al. 2003, PASJ, 55, 473-481
Morii et al. 2016, PASJ, 68, S11
Kong et al. 2010, ATEL #2782

159

Long-term variability of Cen X-3 with MAXIJosé Joaquı́n Rodes,1 Tatehiro Mihara,2 José Miguel Torrejón,1Satoshi Nakahira,3 Guillermo Bernabéu1 and Graciela Sanjurjo1
1

DFESTS, University of Alicante, 03080 Alicante, SpainMAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
3
ISS Science Project Oﬃce, ISAS, JAXA, 2-1-1 Sengen, Tsukuba, Ibaraki 305-8505, Japan
E-mail(JJR): jjrodes@ua.es
2

Abstract
Cen X-3 is an O-type donor Supergiant X-ray Binary with a neutron star as a compact object. The
aim of this work is to study both the light curve and orbital phase spectroscopy of this system in the longterm. Here we estimate the orbital period from the light curve and then we derive the good time interval toextract the spectrum from diﬀerent orbital phases. We obtain orbital phase-averaged and phase-resolved
spectra and we analyse the variability of the model parameters to infer properties of the circumstellar
environment of Cen X-3 and possible dependences of the spectral parameters. The MAXI spectra in the 220 keV energy range were ﬁtted with a model of Comptonisation of cool photons on hot electrons modiﬁedby a partial covering absorption and a ﬂuorescence iron emission line at 6.4 keV.
Key words: X-rays: binaries — pulsars: individual: Cen X-3 — stars: supergiants

The Monitor of All Sky X-ray Image (MAXI ) presents
both all-sky coverage and moderate energy resolution,
which gives us the possibility to investigate the orbital
light curves and the orbital phased-resolved spectra of
Cen X-3 (Rodes-Roca et al. 2017). MAXI observing
strategy suppress eﬀectively the short term variability
associated with accretion and enhances the long term,
permanent structures present in the stellar wind and circumsource environment (Doroshenko et al. 2013; Islamet al. 2014).
2. Timing analysis
First, we have obtained the MAXI/GSC on-demand
light-curve of Cen X-3 from MJD 55 058 to MJD 57 315(i.e. more than 6 years) in the 2–20 keV energy band.
Then, we have searched for a period assuming a sinusoidal signal and derived a Porb = 2.0870 ± 0.0006 days
which is consistent with the value given by Nagase et al.

(1992). In ﬁgure 1 we show the folded light-curve where
we marked the orbital phase range to obtain our orbital
phase-resolved spectra.
0.08
III

V

0.06

counts/bin

1. Introduction
The eclipsing high mass X-ray binary (HMXB) pulsar
Cen X-3 has an O-type supergiant companion, V779
Cen. The orbital period of the binary system is ∼2.1
days with a low eccentricity of ≤0.0016 (Bildsten et al.
1997) and the magnetised neutron star has a spin period of ∼4.8 s (Schreier et al. 1972). The distance to thesource was estimated to be ∼8 kpc (Krzeminski 1974),
although Thompson et al. (2009) gave a value of 5.7±
1.5 kpc.

0.04
I

II

IV

VI

VII

VIII

IX

X

0.02
P = 180315.125000000 s
To = 56185.832782700
0Nbin = 200
í0.4

í0.2

0
orbital phase

0.2

0.4

Fig. 1. Background subtracted orbital phase light curves.

3. Spectral analysis
We have extracted the orbital phase-averaged spectrumof Cen X-3 (see ﬁgure 2) with MAXI/GSC using the
MAXI on-demand processing, carefully excluding any
contamination by nearby brighter sources. We tested
both phenomenological and physical models commonly

160

φ=[0.3í0.4]
normalized counts sí1 keVí1

0.01
φ=[0.6í0.73] (preíingress)10í3

φ=[0.77í0.96] (eclipse)

10í4

2

χ

applied to accreting X-ray pulsars. We rebined all extracted spectra to obtain spectral bins Gaussian distributed. The best ﬁt was attained using CompST, in
Xspec terminology (Sunyaev et al. 1980), modiﬁed by apartial covering model implementation of the TübingenBoulder absorption model (TBnew ) which includes themost up to date absorption cross sections and abundances (Wilms et al. 2000). We needed to put an additional Gaussian absorption line at ∼5 keV to compensate the incompleteness of the GSC response (Nakahira
private communication). Moreover, the ﬂuorescence ironemission line at ∼6.4 keV was also present in the spectrum.

0
í2
5

0.02
0.01

4
3
2
1

Unabsorbed flux
NH (10í22 cmí2)

10 20 30 40

Then we have obtained orbital phase-resolved spectraof the HMXB pulsar Cen X-3 accumulating the 60 s duration scans into 10 orbital phase bins. We have ﬁtted
the orbital phase-resolved spectra with the same model
we used in the orbital phase averaged spectrum. Nevertheless, no absorption Gaussian component at 5 keV wasneeded in these ﬁts. We note that the optical counterpart blocked most of the X-ray emission in eclipse in all
the energy bands, so there is no scattering in the wind.
For the ﬂuorescence iron emission line, a Gaussian proﬁle is assumed. Moreover, this iron line is not detectedboth in the eclipse and egress spectra nor resolved thecomplex iron line in the energy range 6.4–6.7 keV.
The unabsorbed ﬂux appears nearly constant in theorbital phase range [0.1 − 0.4], F(2−20) keV ∼ 4.4 × 109
erg cm−2 s−1 , showing that the accretion rate is quite
stable. The relative ﬂux during the eclipse is a factor of 4
smaller than during out of eclipse. The equivalent width(EW) of the Fe Kα is consistent with a constant valueof ∼ 220 ± 60 eV taking into account the uncertainties.
The centre of the energy of the iron line changes with(in units of 10í9 erg/cm2/s)

300

Fig. 2. Orbital phase averaged spectrum of Cen X-3 in 2-20 keV band
and X-ray model continuum modiﬁed by the iron emission line at
6.4 keV. Top panel: Data and model. Bottom panel shows theresiduals between the spectrum and the model.

200

10Energy (keV)

EWFeKα (eV)5orbital phase suggesting the coexistence of two iron lines
at diﬀerent energies (see ﬁgure 4).
This work was supported by the project ESP201453672-C3-3-P and partially by GV2014/088. TM acknowledges the support by the grant by the Universityof Alicante under visiting research programme INV1508. JJRR acknowledges the support by the Matsumae
International Foundation Research Fellowship No14G04.

100

5×10í3
2×10í30

2

normalized counts sí1 keVí1

Fig. 3. Orbital phase resolved spectra of Cen X-3 in 2-20 keV band.
Top panel: Selected spectra and absorbed CompST plus Fe Kαemission line. Bottom panel shows the residuals for the model.

í2

χ

10
Energy (keV)

0.2

0.4

0.6

0.8

1

Orbital phaseFig. 4. Orbital phase changes in the free parameters, top panel: X-ray
unabsorbed ﬂux, middle panel: column density, and bottom panel:
EW of the Fe Kα.

References
Bildsten, L. et al. 1997, ApJS, 113, 367
Doroshenko, V. et al. A&A, 554, A37
Islam, N. et al. 2014, MNRAS, 441, 2539
Krzeminski, W. 1974, ApJ, 192, L135
Nagase, F. et al. 1992, ApJ, 396, 147
Rodes-Roca et al. 2017, A&A (to be submitted)
Schreier, E. et al. 1972, ApJ, 172, L79
Sunyaev, R. A. et al. 1980, A&A, 86, 121
Thompson, T. W. et al. 2009, Apj, 691, 1744
Wilms, J. et al. 2000, ApJ, 542, 914

161

Orbital resolved spectroscopy of GX 301–2 with MAXI–GSC
Nazma Islam,1 and Biswajit Paul2
1

Nicholas Copernicus Astronomical Center, Warsaw, Poland
2
Raman Research Institute, Bangalore, India
E-mail(NI): nislam@camk.edu.pl
Abstract

The highly absorbed bright high mass X-ray binary GX 301–2, exhibits stable periodic orbital intensity
modulations with a strong pre-periastron X-ray ﬂare. Several models have been proposed to explain theaccretion at diﬀerent orbital phases, invoking accretion via stellar wind or equatorial disc, and accretion
stream from the companion star. In Islam & Paul (2014), we presented results from exhaustive orbital
phase resolved spectroscopic measurements of GX 301–2 using data from the Gas Slit Camera onboard
MAXI. We have found a strong orbital dependence of the absorption column density and equivalent widthof the iron emission line. A very large equivalent width of the iron line along with a small value of thecolumn density in the orbital phase range 0.1–0.3 after the periastron passage indicates the presence of highdensity absorbing matter behind the neutron star in this orbital phase range. The orbital dependence of
these parameters are then used to examine the various accretion models of GX 301–2 and provide stronger
constraints.
Key words: stars: individual: GX 301-2 stars: neutron X-rays: stars

1. Introduction
GX 301-2 is a high-mass X-ray binary (HMXB) systemwith an X-ray pulsar and a B-emission line hypergiant
star WRAY 997. The orbital period of the binary systemis ∼ 41.5 d and a spin period of ∼ 685 s (Koh et al. 1997).
It exhibits periodically varying intensity modulations: abright phase during X-ray ﬂare (pre-periastron passage
around orbital phase 0.95), dim or low intensity phase(after periastron passage around orbital phase 0.15–0.3)
and intermediate intensity phase (during the apastronpassage around orbital phase 0.5). A strong X-ray ﬂareoccurs before the periastron passage as well as a medium
intensity peak is observed at the apastron passage, indicating accretion on to the neutron star due to bothspherical stellar wind along with a possible equatorial
disc or accretion stream (Pravdo & Ghosh 2001; Leahy& Kostka 2008).
GX 301–2 has a highly absorbed X-ray spectrum witha partial covering high energy cutoﬀ power-law component and several emission lines. It has a very high line ofsight photoelectric absorption, which is attributed to thedense circumstellar environment in which the neutronstar moves. The column density varies strongly with orbital phase with certain amount of clumpiness attributed
to the stellar wind (Mukherjee & Paul 2004). A prominent Fe Kα line is found to exist in almost all orbital
phases. This ﬂuorescence line is produced due to repro-

cessing of X-ray photons from the pulsar by the circumstellar matter. The equivalent width of the Fe Kα line
depends on the distribution (geometry and column density) of the surrounding matter (Inoue 1985; Kallman etal. 2004). Therefore, by comparing the equivalent widthof Fe Kα line with NH , we can study the distribution of
circumstellar matter around the neutron star at diﬀerent orbital phases and can be further used to examinevarious accretion models.
In Islam & Paul (2014), we carried out orbital phaseresolved spectroscopic study of GX 301–2, using MAXI–
Gas Slit Camera (Matsuoka et al. 2009). Using spectroscopic analysis of the MAXI data with unprecedented
orbital coverage for many orbits continuously, we studied the orbital phase dependence of the column density
and the line equivalent width, which are then used to
examine the various models about the distribution of
circumstellar matter and the mode of accretion in GX
301–2.
1.1. Data and Analysis
Using MAXI on demand data, 1 we extracted orbital
resolved spectra in 21 independent orbital bins. These
orbital resolved spectra are ﬁtted with two models: anabsorbed power-law continuum, with and without a highenergy cut-oﬀ. A Fe ﬂuorescence line was found in all theorbital phases, which was modelled by a single Gaussian*1 http://maxi.riken.jp/mxondem/

Γ

162

1
0.5
0

1500

Equivalent width (eV)

20
0
0.04
0.02
15000

Flux

1000
500
8
6
4
2

ExcessEqw Line fluxNH

40

8×10í3
6×10í3
4×10í3
2×10í3

1000

500

0

0

0.5

1
Orbital Phase

1.5

10

20
NH (1022 cmí2)

30

40

2

Fig. 1. Left panel (a): Orbital variation of Photon index (Γ), column density (NH in 1022 cm−2 ), Line ﬂux of Fe Kα (photons cm−2 s−1 ),Equivalent width of Fe Kα line (Eqw in eV) and Flux of source (F in 10−9 ergs s−1 cm−2 ) for power-law model with high energy cut-oﬀ
model. The arrow denotes the orbital phase of high equivalent width of Fe Kα line for low NH . Circles corresponds to the spectral ﬁtvalue taken from Endo et al. (2002) (ASCA), squares from La Barbera et al. (2005) (BeppoSAX), triangles from Suchy et al. (2012)(Suzaku) and star from Fürst et al. (2011) (XMM-Newton). Right panel (b): Plot of equivalent width of Fe Kα versus NH . Solid line anddashed lines represents the relation between equivalent width and column density of absorbing matter for isotropically distributed matterat diﬀerent Γ (Inoue 1985; Kallman et al. 2004).

line. For some orbital phases near the X-ray peak, a low
energy excess is found to be present in the spectra. To
only estimate the ﬂux in the soft excess, we have modelled the low energy excess with an unabsorbed blackbody component.
2. Discussions and Conclusions
Figure 1a shows the orbital variation of Γ, NH , ﬂux andequivalent width of Fe ﬂuoresence line, total ﬂux of thesystem and ratio of ﬂux included in the low excess to thetotal ﬂux, for the absorbed power-law with a high energy
cutoﬀ model. The NH and equivalent width measurements of the Fe Kα using single pointed observations at
diﬀerent orbital phases are overlaid on the plot and theyshow a large scatter in the measurements. The shortscale variations are smeared out while studying the longterm averaged behaviour of this source with MAXI–GSC
data and the long term accretion structures are brought
forth.
The column density NH is found to vary with a pattern similar to the ﬂux of the system, indicating a possible origin of ﬂare due to increased mass accretion. The
orbital variation of equivalent width of Fe Kα line showsa diﬀerent trend as compared to the orbital variation ofcolumn density. The highest equivalent width occurs at
the dim phase of 0.1–0.3 which also has lowest NH along
the line of sight (denoted by arrows in Figure 1a). Figure 1b is the plot of equivalent width of the iron line
and the absorbing column density NH in diﬀerent orbital bins. These observations highly deviate from therelation expected for an isotropically distributed gas (Inoue 1985; Kallman et al. 2004). Instead, there seems toexist high anisotropicity in the distribution of circumstel-lar matter around the X-ray pulsar, especially in someorbital phases. These results strongly favour a high density gas stream plus a stellar wind model for mode ofaccretion on to the neutron star in GX 301–2 and provide stronger constraints to the model (Leahy & Kostka
2008). The relation between equivalent width of Fe line
and column density of absorbing matter can be used forprobing the geometry and distribution of circumstellar
matter around other wind-fed systems.
References
Endo, T., Ishida, M., Masai, K., et al. 2002, ApJ, 574,
879
Fürst, F., Suchy, S., Kreykenbohm, I., et al. 2011, A&A,
535, A9
Inoue, H. 1985, SSRv, 40, 317
Islam, N., & Paul, B. 2014, MNRAS, 441, 2539
Kallman, T. R., Palmeri, P., Bautista, M. A., Mendoza,
C., & Krolik, J. H. 2004, ApJS, 155, 675
Koh, D. T., Bildsten, L., Chakrabarty, D., et al. 1997,
ApJ, 479, 933.
La Barbera, A., Segreto, A., Santangelo, A.,
Kreykenbohm, I., & Orlandini, M. 2005, A&A, 438,
617
Leahy, D. A., & Kostka, M. 2008, MNRAS, 384, 747
Matsuoka, M., Kawasaki, K., Ueno, S., et al. 2009, PASJ,61, 999
Mukherjee, U., & Paul, B. 2004, A&A, 427, 567
Pravdo, S. H., & Ghosh, P. 2001, ApJ, 554, 383
Suchy, S., Fürst, F., Pottschmidt, K., et al. 2012, ApJ,
745, 124
Watanabe, S., Sako, M., Ishida, M., et al. 2003, ApJL,
597, L37

163

SFXTs vs sgHMXBs: does the diﬀerence lie in the companion wind?
Pragati Pradhan,1 Biswajit Paul,2 Enrico Bozzo,31

St. Joseph’s College, Singamari, Darjeeling-734104, West Bengal, India
2
Raman Research Institute, Sadashivnagar, Bangalore-560080, India
3
ISDC, University of Geneva, Chemin d’Ecogia 16, Versoix, 1290, Switzerland
E-mail(PP): pragati2707@gmail.com
Abstract
We present a comparative study of Supergiant X-ray binary (sgHMXBs) and Supergiant Fast X-ray
Transient (SFXTs) systems by using the absorption column density and equivalent width of iron Kα line intheir X-ray emission. The work has been carried out using 71(33) out-of-eclipse observations of sgHMXBs(SFXTs) with Suzaku and XMM-Newton and we have taken care to separately analyse parts of any
observation with signiﬁcant variation in the spectrum. Analysis of all archival Suzaku and XMM-Newton
observations of these systems show that sgHMXBs have a wide range of equivalent width of iron emission
line and equivalent column density of absorption, both over three orders of magnitude. In comparison,
the SFXTs show a smaller range for both the parameters, less than one and a half order of magnitude.
These ﬁndings indicate a crucial diﬀerence in the wind characteristics of the companions of
sgHMXBs and SFXTs, which could be an important factor for the intriguing diﬀerence inaverage X-ray luminosity and transient behaviour between these two classes of sources.
Key words: X-rays: binaries–pulsars: general

1. Introduction
SFXTs (Sguera et al., 2005) are a subset of the sgHMXBs that are characterized by short outbursts with
fast rise times (∼ tens of minutes) and typical durationsof a few hours. Except during these bursts, SFXTs have
X-ray luminosity that is 2-4 orders of magnitude smaller
than the classical supergiant HMXBs. The compact object is believed to be a neutron star accreting from thestellar wind of an O or B-type supergiant. However,
X-ray pulsations have been conﬁrmed from only a few
of the SFXTs (eg, IGR J00370+6122, Grunhut, Bolton& McSwain 2014) and cyclotron line has been detected
in only two SFXTs (IGR J17544-2619; Bhalerao et al.
2015; Bozzo et al. 2016 and IGR J18483-0311; Sguera
et al. 2010), that awaits reconﬁrmation. SFXTs, havefor long been speculated to be descendants of Be X-ray
binaries, more so because some of these systems show
evidence of disk-like structures (Ducci, Sidoli & Paizis,2010) around the supergiant similar to Be/X-ray binaries. The X-ray behaviour are diﬀerent in spite of binary parameters and the stellar companions being quite
similar in sgHMXBs and SFXTs. There is, therefore, agrowing interest to make a comparative study of classical HMXBs with SFXTs to understand the origin of
their diﬀerent X-ray characteristics of such systems.

In this work, we make the ﬁrst of it’s kind study of thewind of the companion star in the two class of sources
through the absorbing column density of material in theline of sight and the iron ﬂuorescence line to understandthe diﬀerences/similarities in these two class of HMXBs
as a whole.
2. Observations and data reduction
For the purpose of analysis, we have taken up archived
data (till July 2016) for both classical HMXBs andSFXTs observed with Suzaku (Mitsuda et al., 2007) andXMM (Jansen et al., 2001) and extracted spectrum asfollows:• Suzaku (XMM) data: Extracted XIS0 (PN/MOS)
spectrum following procedure in Suzaku guide 1(XMM analysis thread 2 ).
• X-ray spectra (0.3-10.0 keV) for each observation
were extracted from time segment with unvarying
spectral parameters and ﬁtted with a spectral model
consisting of an absorbed power-law and a gaussian
emission line.
*1 http://heasarc.gsfc.nasa.gov/docs/suzaku/analysis/abc/*2 http://www.cosmos.esa.int/web/xmm-newton/sas-threads

164

ferent accretion regimes that cause change in X-ray
luminosity. Such gating mechanisms, however, require large periods and magnetar-like magnetic ﬁelds (see eg, Bozzo, Falanga & Stella 2008).

NH(1022 atoms cmí2)

100
SFXTs
4U 1907+09
Cen Xí3
EXO 1722í363
GX 301í2
IGR J16318í4848
IGR J16320í4751
IGR J16393í4641
IGR J16393í4643
LMC Xí4
4U 1700í37
IGR J16493í4348
4U 1909+07
OAO 1657í415
4U 0114+65
IGR J16207í5129
4U 1538í522
SMC Xí1
Vela Xí1
SAX J1802.7í2017

10

1

0.1

0.01

0.1

1

10

EQW (keV)
Fig. 1. Plot of the column density versus equivalent width of FeKα line. SFXTs are marked in black while diﬀerent sgHMXBs
are labelled.

3. Results
Figure 1 shows a plot of the EQW of Kα line with equivalent column density of hydrogen (NH ) obtained from all
the Suzaku and XMM observations.

• Clumpy winds: Inhomogeneties of circumstellar
matter around neutron star cause varied accretion,and hence varied luminosity. Useful to explain
only short ﬂare SFXTs like in the case of IGR
J17544-2619 (Rampy, Smith & Negueruela, 2009).
It is clear from the above discussions that the cause
for diﬀerence in behaviour between sgHMXBs andSFXTs are many but not unique. In this work, we
propose one possible explanation for the diﬀerence intheir behaviour. We ﬁnd that sgHMXBs are in generalmore absorbed than SFXTs. This could be possible
if the neutron star orbits a denser medium in
sgHMXBs because of slower stellar wind of thedonor. This also explains large equivalent width of theemission line in some sgHMXBs. The wind terminal
velocity seem to have a decisive role in determining the
behaviour of these two class of HMXBs. This is in agreement with the study of wind properties of two membersthat act as prototypes for SFXTs (IGR J17544-2619)
and sgHMXBs (Vela X-1) where it was found that
their behavioural diﬀerence is because of diﬀerences intheir wind properties only (Giménez-Garcı́a et al., 2016).

References

4. Discussion
We ﬁnd that for SFXTs, the iron line has smaller equivalent width compared to all other sgHMXBs. The column
density of SFXTs have a smaller spread of 1.5 decades
compared to the sgHMXBs which have a spread of overthree decades. This shows a crucial diﬀerence in theenvironment surrounding the neutron star in these twoclasses.
Even though sgHMXBs and SFXTs are similar with regards to the compact object, companion star and theorbital period, their dynamic range of X-ray luminosities is very much varied. Some prominent theories put
forward so far to explain such a behaviour are:• Geometrical diﬀerences: High eccentricity of
SFXTs and nearly circular orbits of sgHMXBs cause
sudden ﬂaring in SFXTs and persistent emission in
sgHMXBs. This explanation is, however, put into
question by the discovery of SFXTs of shortorbital periods like IGR J18483-0311 and IGR
J16479-4514 (Jain, Paul & Dutta, 2009a,b).
• Centrifugal/gating mechanisms: Onset of dif-

Bhalerao V. et al., 2015, MNRAS, 447, 2274
Bozzo E. et al., 2016, A&A, 596, A16
Bozzo E., Falanga M., Stella L., 2008, ApJ, 683, 1031
Ducci L., Sidoli L., Paizis A., 2010, MNRAS, 408, 1540
Giménez-Garcı́a A. et al., 2016, A&A, 591, A26
Grunhut J. H., Bolton C. T., McSwain M. V., 2014, A&A,
563, A1
Jain C., Paul B., Dutta A., 2009a, MNRAS, 397, L11
Jain C., Paul B., Dutta A., 2009b, Research in Astronomy
and Astrophysics, 9, 1303
Jansen F. et al., 2001, A&A, 365, L1
Mitsuda K. et al., 2007, PASJ, 59, 1
Rampy R. A., Smith D. M., Negueruela I., 2009, ApJ, 707,
243
Sguera V. et al., 2005, A&A, 444, 221
Sguera V., Ducci L., Sidoli L., Bazzano A., Bassani L., 2010,
MNRAS, 402, L49

165

Indications for a long-term periodicity in EXO 2030+375
Eva Laplace 1,2 , Tatehiro Mihara 2 , Yuki Moritani 3 , Motoki Nakajima 4 ,Toshihiro Takagi 2 , Kazuo Makishima 2 and Andrea Santangelo 1
1

Institut für Astronomie und Astrophysik, University of Tübingen, Sand 1, 72076 Tübingen, Germany
2
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
3
Kavli Institute for the Physics and Mathematics of the Universe (WPI),The University of Tokyo Institutes for Advanced Study, The University of Tokyo, Kashiwa, Chiba 277-8583, Japan
4
School of Dentistry at Matsudo, Nihon University, 2-870-1, Sakaecho-nishi, Matsudo, Chiba 271-8587, Japan
E-mail(EL): laplace@astro.uni-tuebingen.de
Abstract
The X-ray source EXO 2030+375 (spin period: 42 s; orbital period: 46.021 d) is known for being the
Be X-ray binary system showing the largest number of recurring low-luminosity X-ray outbursts (type I)
every orbital period. Recent changes in the source behavior, which include a fading of the X-ray ﬂux anda torque reversal in 2016, were found to be very similar to events which occurred 20.5yr earlier, just before
the source experienced a sudden shift of the orbital phase of the outburst peak (orbital phase jump) in
1995 (Wilson et al. 2002). Moreover, the observation that the type II outbursts observed in 1985 (Parmar
et al. 1989) and 2006 (Corbet et al. 2006) occurred at times just in-between these events lead to theinterpretation that the source has a cyclic behavior with a time-scale of 21 or 10.5 yr. In our previous
paper (Laplace et al. 2017), we predicted the future behavior of the source depending on two possible
interpretation. Here, we report the observation of an orbital phase shift in July 2016 which conﬁrms theprediction and supports the existence of a 21 yr periodicity of the source behavior, which is best explainedby Kozai-Lidov oscillations of the circumstellar disk.
Key words: stars: individual: EXO 2030+375 –stars: neutron, emission-line, Be – X-rays: binaries

1. Introduction
EXO 2030+375 was discovered in 1985 (Parmar et al.
1989) and is characterized by an orbital period of 46.02
d and a spin period of 42 s (Wilson et al. 2002). It
is composed by a neutron star and a BO Ve companion(Coe et al. 1988) and as such belongs to the class of
Be/X-ray binaries. These sources are characterized bya transient X-ray radiation, which is divided into two
characteristic types: (i) type I outburst, which typicallyoccur every orbital period and are understood as the accretion of matter from the companion star at periastron,(ii) type II outbursts, very luminous outbursts which canlast for several orbital periods. Their origin remains unclear, with accretion of large amounts of matter from awarped and eccentric disk surrounding the Be star themost likely candidate (Okazaki et al. 2014; Martin etal. 2014a).
2. Recurrent behavior
In Laplace et al. (2017), we investigated the relationship between recent changes in the source behavior andsimilar events in 1995. The source had recently shown anunusual drop in X-ray ﬂux and a change from a steadily
increasing spin frequency to a constant spin frequency(Fuerst et al.
2016). We noticed that the type II
outbursts observed in 1985 (Parmar et al. 1985) and2006 (Corbet & Levine 2006) were occurring at peculiar
times, just in between the drops in ﬂux and spin frequency derivative. The data suggested a recurrence ofabout 21 years with alternating particular events (orbital
phase shifts and type II outbursts) every 10.5 years.
After studying possible interpretations for this phenomenon, we found that the time-scale of Kozai-Lidov
oscillations predicted for EXO 2030+375 nicely matcheda 10.5 period between negative orbital phase shifts andgiant outbursts. Based on this interpretation, we predicted the observation of a giant outburst or of an orbital phase shift around 2016 December, depending on
the Be disk state (highly eccentric or highly inclined withrespect to the orbital plane). Here, we present observations of an orbital phase shift in June 2016 (ﬁrst reported
in Laplace et al. (2016)), which conﬁrm our prediction
and indicate that there is indeed a recurrent pattern inthe source.

166

0.15

Flux [Crab]

0.10

Swift/BAT orbit 155
Swift/BAT orbit 164
MAXI/GSC orbit 155
MAXI/GSC orbit 164

0.05
0.00

−0.05
−0.4

−0.2

0.0

0.2

Orbital phase

0.4

0.6

Fig. 1. Folded MAXI/GSC and Swift/BAT light-curve of EXO 2030+375. See text for more information.

3. Observation of an orbital phase shift
Data for our observations were obtained from
MAXI/GSC and Swift/BAT, which are publicly
available. The outburst occurred earlier than expected.
We folded the light-curve with the orbital period andobserved a shift of the outburst peak of about 4.3 days,
as shown in Fig. 1. The peak of the outburst was
attained on 2016 July 20, 5 days earlier than the almost
constant peak time of the last 7 years. In terms of theorbital phase, the peak of the outburst shifted from an
orbital phase of 0.13 to 0.015. It is interesting to noticethat the phase shift was in the same direction as theone in 1995, that is, earlier than the previous peak time.
The orbits are calculated starting from the time of theﬁrst periastron passage of the source observed by the
RXTE/ASM sattelite, MJD 50086.967.
3.1. Discussion and Conclusion
In our previous work, we predicted the observation of anorbital phase shift or of a type II outburst, depending on
the interpretation used for the recurrent behavior. The
observation of an orbital phase shift was predicted in thecase of recurrent Kozai-Lidov oscillations in the circumstellar disk and the observation we report here conﬁrms
this prediction. However, it does not constrain the oscillation time-scale discussed in our previous work, sincesuch an observation was expected for both an eccentric,
low-density disk and a highly inclined disk.
The fact that the orbital phase shift occurred much earlier than what we expected can be explained by the large
uncertainty of our estimation and by the fact that simulations of Kozai-Lidov oscillations in hydrodynamicaldisks (Martin et al. 2014b; Fu et al. 2015) have shownthat the oscillations can vary depending on various parameters, including the disk viscosity, the disk warp andthe disk tilt, leading to a diﬀerent time-scale of the oscillations. However, the time-scale of Kozai-Lidov oscillations calculated in our previous work suggest that orbital
phase shifts occur when the disk is highly eccentric. Optical and infrared observations are needed to verify thedisk state. To obtain and analyze these data constitutes
our next challenge.
References
Coe, M. J. et al. 1988, MNRAS, 232, 865
Corbet, R. H. D. et al. 2006, ATel, 843
Fu, W., et al. 2015, ApJ, 807, 75
Fuerst, F. et al. 2016, ATel, 8835
Krimm, H. A. et al. 2013, ApJS, 209, 14
Laplace, E. et al. 2016, ATel, 9263
Laplace, E. et al. 2017, A&A, 597, A124
Martin, R. G. et al. 2014a, ApJ, 790, L34
Martin, R. G. et al. 2014b, ApJ, 792, L33
Matsuoka, M. et al. 2009, PASJ, 61, 999
Okazaki, A. T., & Negueruela, I. 2001, A&A, 377, 161
Okazaki, A. T. et al. 2013, PASJ, 65, 41
Parmar, A. N. et al. 1989, ApJ, 338, 359
Reig, P. 2011, Astrophys. Space Sci., 332, 1
Wilson, C. A. et al. 2002, ApJ, 570, 287
Wilson, C. A. et al. 2008, ApJ, 678, 1263

167

MAXI Observations of Supergiant Fast X-ray Transients
Ai Sakamaki and Hitoshi Negoro
Department of Physics, Nihon University
1-8 Kanda-Surugadai, Chiyoda-ku, Tokyo, Japan
E-mail(AS): sakamaki@cygnus.phys.cst.nihon-u.ac.jp
Abstract
About 20 supergiant fast X-ray transients (SFXTs) have been discovered mainly by INTEGRAL. The
MAXI nova-alert system was often triggered by short X-ray ﬂares or outbursts probably from SFXTs,
for instance, IGR J18483−0311, AX J1841.0−0536, and AX 1739.1−3020 (aka XTE J1739−302/IGR
J17391−3021) even though about half of SFXTs can not be spatially resolved from nearby bright X-ray
sources. Thus, MAXI may discover new SFXTs in future (a present SFXT candidate is MAXI J1932+091).
Here, we brieﬂy summarize preliminary MAXI results of SFXTs, and present the capability to detect theorbital periodicity of IGR J18483−0311.
Key words: X-rays: transient — neutron star: SFXTs

1. Introduction
Supergaint Fast X-ray Transients, SFXTs, (e.g.,Sguera
et al. 2005; Negueruela et al. 2006) are a sub-class ofhigh mass X-ray binaries (HMXBs). SFXTs exhibit relatively short, and not-so-bright hard X-ray ﬂares. The
ﬂare activities typically last less than one hour, and thepeak ﬂux of the ﬂares are around 100 mCrab or less inthe soft X-ray band, which are just below the detection
limits of the previous all sky monitors, e.g., Gigna/ASM
and RXTE/ASM. These characteristics of the sources
had long hidden their nature until INTEGRAL discovered them.
The MAXI nova-alert system (Negoro et al. 2016) was
often triggered by probable short transient activities of
SFXTs in a single scan transit or 4 orbits (∼ 6 h) timebin. If detected, we, the MAXI team, provide a prompt
E-mail alert to the members of the x-ray-star MAXI
mailing list. Here, we brieﬂy summarize observational
properties of SFXTs from MAXI/GSC observations.
2. MAXI detections of SFXTs
About half of SFXTs are so close to bright sources thatMAXI can not resolve the sources. Such sources areshown in light green in ﬁgure 1 (see a color version on thewebsite). Furthermore, X-ray properties of ﬂares from
SFXTs, ”short and hard”, make it diﬃcult for MAXI tocatch ﬂares. Nevertheless, X-ray enhancements probably due to activities of SFXTs often triggered the MAXI
nova-alert system, or were recognized on GSC images,e.g., AX J1841.0−0536 (Negoro et al. 2010) and XTE
J1739−302/IGR 17391−3021 (Negoro et al. 2015). Usu-

ally, such enhancements were recognized only in a single
scan orbit, except for outbursts from IGR J18483−0311.
We, however, could not conﬁrm that the activities werereally due to SFXTs except for simultaneous detection
with other observatories.
Such possible detections of SFXTs with MAXI/GSC
are summarize in table 1. As shown in the table, thenumbers of the MAXI detections are positively correlated with those of Swift/BAT detections, suggestingthat MAXI actually detected the source activities. We
are checking the signiﬁcance of each detection and thesource location more precisely by taking point spreadfunctions of the cameras and nearby sources into account.
3. Periodicity
Periodicity search was performed for IGR J18483−0311
using GSC data over 7 years, and the orbital periodof ∼ 18.5 days was clearly detected in the power spectrum density (ﬁgure 2). Sakakibara (2013) obtained amore precise period of 18.576 ± 0.016 days using 4 years
MAXI/GSC data.
He also showed that a 2–20 keV folded light curveof GSC data had a proﬁle very similar to a 15–50 keVlight curve of Swift/BAT data, and that the soft X-rays
tended to increase more gradually than the hard X-rays.
Such gradual increase in the soft X-ray band was also
recognized in RXTE/ASM data (Levine, et al. 2011).
4. Discussion and Future Works
As shown above, MAXI can detect SFXT activities.
This implies that undiscovered SFXTs will be detected

168

GX 3+1
IGR J18462-0223
AX J1749.1-27333
2X
XMM
M J185114.3-000004
IGR J19294+1816

AX J1841.0-053
5336

XTE J1901+01
0114

IGR
IG
R J1
J174
7407
07-2808

IGR J16418-4532
AX J1739
9.1.1-3020
1
4U 1630-472

GX 13+1

SAX J1818.6-1703

IGR J184833-0
-031
03 1
IGR
R J1
J18175-2241
419
19
MAXI J1932+091
AX J1845.0-0433

IGR J112215-5952

IG
GR J166207-5
512
129

IGR
R J1
16465
6465
6 -4
- 50
07
IGR J16464
479
79-45147951
14IG
GR J1175
7 44-261
44 61
44
6199

IG
GR J1
J161
61995-494
45

Cen X-3

IGR J1
173
7354
5 -3
54
-325
2 5

IGR
R J0
J0840808-4-4503

Vela Pulsar

IGR
R J1
J166328-472
72
26

GX 340+0

Fig. 1. Locations of known SFXTs and nearby bright sources shown on a MAXI GSC+SSC X-ray image.

Table 1. List of SFXTs spatially resolved by MAXI (shown in yellow in ﬁgure 1)(264.041, −1.950)(332.459, −1.050)

MAXI
Detection1
4
2?

BAT
Detection1
10(358.068, +0.445)

8?

27(26.764, −0.239)

10

26(28.138, −0.658)

3?

(29.750, −0.745)(32.909, −0.043)(35.381, −1.623)(53.540, +0.115)

Object Name(l, b)IGR J08408−4503
IGR J16207−5129
AX J1739.1−3020/XTE J1739−302/IGR J17391−3021
AX J1841.0−0536
AX J1845.0−0433/
IGR J18450−0435
IGR J18483−0311(+ IGR J18462−0223)
2XMM J185114.3
XTE J1901+014
IGR J19294+1816

Nearby SourceName
f 2 [mCrab]
Vela Pulsar
∼ 10

Δθ[d]
0.54

IGR J1735−3022
XB 1732−304

∼20

0.35
0.74

14

Swift J1843.5−0343

∼ 10

0.86

35

57

GS 1843−02

∼ 10

0.75
(0.38)

0
3
2?

2

1. All the numbers of the MAXI detections are preliminary. Those of the BAT detections are from Romano et al. (2014) except
for XTE J1901+014 (Krimm et al. 2010). The numbers show total days, not the times of triggers or outbursts.
2. Steady 4-10 keV X-ray ﬂux, including the galactic ridge X-ray emission, observed with GSC at the nearby source region.

1.5×10í310í3
5×10í40intensity [arbitrary]

2×10í3

not be excluded.
On the other hand, we have to carefully reanalyze
MAXI data of known SFXTs as described above to evaluate long-term source activities or its recurrence in thesoft X-ray band, which have not been extensively investigated after e.g., Levine, et al. (2011).

02×10í7

4×10í7
6×10í7frequency [Hz]8×10í7

10í6

Fig. 2. PSD of IGR J18483−0311 (MJD55058-57692, 4.0-10.0 keV)in future. In fact, MAXI J1932+091 (Yamaoka et al.
2014; Kennea et al. 2014), of which the companion is
likely to be a Be star (Itoh et al. 2014), is a SFXT candidate though it locates at relatively high galactic latitude,b ∼ −5 deg (see ﬁgure 1), and its emission lasted morethan 9 hours (also see Negoro et al. 2016). Of course,
other possibilities, for instance, a black hole binary, canReferences
Itoh, A. et al. 2014, Astronomer’s Telegram 6186
Kennea, J. et al. 2014, Astronomer’s Telegram 6177
Krimm, H. et al. 2010, Astronomer’s Telegram 2375
Levine, A. et al. 2011, ApJS, 196, 6
Matsuoka, M. et al. 2009, PASJ, 61, 999
Negoro, H. et al. 2010, Astronomer’s Telegram 3018
Negoro, H. et al. 2015, Astronomer’s Telegram 6900
Negoro, H. et al. 2016, PASJ, 68, S1
Negueruela, I. et al. 2006, ApJ, 638, 982
Romano, P. et al. 2014, A&A, 564, A2
Sakakibara, H. 2013, Master’s thesis, Nihon university
Sguera, V. et al. 2005, A&A, 444, 221
Yamaoka, K. et al. 2014, Astronomer’s Telegram 6174

169

Monitoring Supergiant Fast X-ray Transients with Swift and
XMM-Newton
Enrico Bozzo1
1

ISDC - University of Geneva
E-mail(EB): enrico.bozzo@unige.ch
Abstract
In this contribution, I will present the current status of our long term monitoring program of theSupergiant Fast X-ray transients with Swift/XRT, including an exhaustive summary of the main ﬁndingson this class of extremely variable sources in the X-ray domain. I will also report on the most recent results
obtained through a combined XMM and NuSTAR long observational campaign on the SFXT prototype
IGRJ17544−2619, highlighting all challenges that do not currently allow us to theoretically interpret the
SFXT phenomenology in a satisﬁable way.

170

171

Discovery of keV excess emission in the isolated neutron star RX
J1856.5−3754
Tomokage Yoneyama1
1

Osaka University Graduate School of ScienceE-mail(TY): yoneyama@ess.sci.osaka-u.ac.jp
Abstract

RX J1856.5−3754 is the brightest and nearest (∼120 pc) thermally emitting isolated neutron star, a
prototype of X-ray dim neutron stars (XDINS). After its discovery with ROSAT, the source was observed
many times with various X-ray satellites. Observations with XMM-Newton and Chandra satellites indicatethat the X-ray spectrum is well reproduced with a combination of blackbody models, one with kT ∼ 32 eV
and the other with kT ∼ 63 eV. In addition, the X-ray spectrum and intensity are found to be fairly stable
over the time. Therefore, the source was employed as a calibration target of the soft X-ray detectors,
which are sensitive to contamination of material onto detectors or ﬁlters inside the satellites. Suzaku
observed this source 10 times during its lifetime primarily for the calibration purpose. The X-ray spectra
below 0.8 keV are well ﬁtted with the two blackbody model determined with XMM-Newton and Chandra.
Nevertheless, we notice systematic excess of the observed spectra over the two blackbody model. The
excess is about 20% in the integrated counts in 0.8–1.2 keV band, and is found both in the spectra of
BI-CCD (XIS1) and FI-CCD(XIS0+XIS3). We further analyze the XMM-Newton EPIC-PN spectra ofthis source and ﬁnd similar excess. We call this keV excess, discovered 1st time for this source and for the
XDINS class. We examine possible causes of this excess, including uncertainty in background subtraction,
pileup of photons below 0.8 keV, and contamination of other sources. We show qualitative diﬃculties in
explaining the keV excess with these causes. We then conclude the keV excess is originated in the source
RX J1856.5−3754, or its neighbor within the radius of 0.2 arcmin. The keV excess component is ﬁtted
either with a blackbody model ( kT ∼ 140 eV) or a power law model (Γ ∼ 4.5), though it is hard toconstrain considering the systematic uncertainty. We shortly discuss possible origin of this keV excess.

172

173

Stellar
Flares174

175

Stellar ﬂares detected with MAXI for the past 7 years
Yohko Tsuboi1 , Ryo Sasaki1 , Yumiko Nakamura1 , Ken Yabuki1 , Atsushi Kawagoe1 ,Soichiro Kaneto1 , Yasuharu Sugawara2 , Hitoshi Negoro3 , and the MAXI Team1

Department of Physics, Faculty of Science and Engineering, Chuo University,
1-13-27 Kasuga, Bunkyo-ku, Tokyo 112-8551, Japan
2
ISAS, JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 229-8510, Japan3
Department of Physics, Nihon University, 1-8 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8308, Japan
E-mail(YT): tsuboi@phys.chuo-u.ac.jp
Abstract
Since the launch in 2009 August, with the unprecedentedly high sensitivity as an all-sky X-ray monitor,
MAXI has caught more than a hundred of huge ﬂares from stars. Most of them are from low-mass, active
stars (RS CVn systems, an Algol system, dMe systems, a dKe system, Young Stellar Objects). With thetotal radiative energy of 1034−39 ergs, the MAXI detections have broken the record of the largest ﬂaring
magnitudes in each stellar categories (e.g. “RS CVn” and so on). The enlarged sample of intense ﬂares
has enabled us to do systematic studies in various viewpoints. One of the studies is our discovery of auniversal correlation between the ﬂare duration and the intrinsic X-ray luminosity, which holds for 5 and
12 orders of magnitude in the duration and LX , respectively (Tsuboi et al. 2016).
Besides low-mass stars, a historically brightest X-ray ﬂare has been detected with MAXI from themassive star system, Eta Carinae (Negoro et al. 2014).
In this review, we introduce the studies of stellar ﬂares obtained in 7.25-year monitoring with MAXI.
Key words: stars: activity — stars: ﬂare — stars: late-type — stars: rotation — stars: variables:
general1.

Introduction

Stellar ﬂares are thought to be a resultant of magnetic reconnection on a stellar surface. The closest stellar ﬂares
to us are solar ﬂares. The typical total energy of a solar ﬂare ranges from 1029 to 1032 erg (e.g., Shibata &Yokoyama 2002), and the duration is several minutes
to several hours (e.g., Shimizu 1995). Despite the long
history of the observation of the solar ﬂares, the highenergy cut-oﬀ of the solar ﬂares in the occurrence rate
(dN/dE) versus ﬂare energy (E) has not yet been reported. Whereas, in 2012, Maehara and his co-workers
discovered 365 ﬂares with the total energy of 1033 –1036
erg, in the range of so called “super-ﬂare” (Schaefer et al.
2000), on 148 G-type dwarfs which includes 10 Sun-likestars (Maehara et al. 2012). This frightens us, because
our civilization would suﬀer from much more severe damages, if superﬂares were really to occur on the Sun.
Even if we are apart from the solar ﬂares, generally,
we cannot answer the following fundamental questions;
how large ﬂares can a star have, and how is such a ﬂare?
The best way to answer such questions would be to use
MAXI, which has the unprecedentedly high sensitivity

for an all-sky X-ray monitor. Here, we report the resultswith the gas proportional counters (GSC) of MAXI obtained in the ﬁrst 7.25-year operation from 2009 August
to 2016 November.
2. Results
2.1. Flare sample
We detected 106 ﬂares on twenty-seven low-mass stars;
fourteen RS CVn systems, one Algol system (Algol),nine dMe stars, one dKe star, one Young Stellar Object(TWA-7) and one K-type variable star. The detectionof the ﬂare from TWA-7 and the results from the ﬁrst
two-year monitoring with MAXI have been already reported in Uzawa et al. (2011) and Tsuboi et al. (2016),
respectively. Most of the categories are classiﬁed as active binaries, which consist of the pair of a sub-giant
and a main sequence, and a pair of a giants and mainsequence, and a pair of two main-sequence stars. The
distance between each stellar component are 1–2 timesof their own stellar sizes. We should note that our sample contains single active stars, although they are minor(Uzawa et al. 2011 and Tsuboi et al. 2016).

176    

       

Fig. 1. Log-log plot of X-ray luminosity in the 2–20 keV band of ﬂares
vs. distance from stars detected with MAXI/GSC. The squares,
triangles, circle, diamond, and star show RS-CVn type stars, dMestars, dKe star, Algol, and TWA-7, respectively. The detection
limit appears to be roughly 10 mCrab in the 2–20 keV band.

Figure 1 shows the X-ray luminosity vs. distance tothe source. This shows that the MAXI sources are within
about 200 pc, and our detection limit is roughly 10
mCrab. The largest ﬂares that we detected in seven-year
monitoring have the luminosity of 1034 ergs s−1 . Thanks
to the improvement of the alert system “nova search”(Negoro et al. 2016) during 7.25 years, all dMe stars
within 10 pc distance have been detected with MAXI.
On the other hand, the detected sample of active binaries are only the tip of the iceberg; within 100 pc distance, 256 active binaries are known to exist, but only
1/10 are detected. Moreover, none of the solar type stars
has been detected, despite that within 20 pc distances,
there are 15 such stars.
The total energies of the detected ﬂares range from1034 to 1039 erg in the 2-20 keV band, i.e. at least, morethan two orders of magnitude larger than the maximum
solar ﬂare (1032 erg). The detail is reported in Sasaki etal. (2017).
2.2. Universal correlations
Figure 2 shows the plot of emission measure vs. plasma
temperature. It had been known that there is a universal
correlation which holds through many orders of magnitudes in emission measure (Shibata & Yokoyama 1999).
MAXI sources lie on the high-end of the correlation. The
large emission measures in the MAXI sources indicatethat the radiative volume, and then the loop size, arelarger than those in the solar ﬂares by many orders ofmagnitudes, and the magnetic ﬁeld, which is indicativeof the plasma temperature, is on the other hand not so
diﬀerent from those in solar ﬂares. In some ﬂares, theFig. 2. Log-log plot of emission measure vs. plasma temperature(kT ) for the MAXI X-ray ﬂares (ﬁlled symbols), along with stellar
ﬂares from RS-CVn type, Algol, dMe stars and YSOs, solar ﬂares,
and solar microﬂares. The arrows indicate the lower limits forindividual MAXI/GSC sources. EM -kT relations, for B= 15, 50,
and 150 Gauss and for the loop-sizes of 108 , 1010 , 1012 , and 1014
cm, from Shibata & Yokoyama (1999) are also indicated.

loop size should be more than 10 times larger than eventhe distance between each binary component.
Recently, Tsuboi et al. (2016) discovered another universal correlation between the intrinsic X-ray luminosity
in the 0.1–100 keV band and the duration of ﬂares. Figure 3 shows the correlation which holds from solar micro
ﬂares to the large ﬂares detected with MAXI/GSC. The
MAXI ﬂares are on the high-end of the positive correlation.
2.3. Serendipitous ﬂares from HD347929
We have detected the ﬂares from HD347929 in 2013
serendipitously and also have recognized that the source
had another ﬂare in 2010. Since we had possessed almost
no information of this source, we conducted a spectroscopic observation in the optical band with Gumma observatory, and obtained a hint for the nature; it seems to
be a young stellar object with a rich Li abundance. The
detail will be reported elsewhere (Tsuboi et al. 2017).
2.4. The historically brightest X-ray ﬂare on the massive
star system, Eta Carinae
MAXI detected not only the ﬂares on the low-mass stars
but also those on a massive stellar system, Eta Carinae.
Eta Carinae is a mystery source with surrounding dust
nebulae. The dust erupted from Eta Carinae itself inthe 19th century. There is also an X-ray emitting nebula, surrounding Eta Carinae, which is located outside ofthe dust nebulae. Then the nature of the central source
has been little known, except for the information thatit seems to be a stellar system containing at least two177 






 


Fig. 3. Log-log plot of duration of ﬂares vs. X-ray luminosity in the 0.1–100 keV band. The best-ﬁt model is inserted with a broad solid line
(Tsuboi et al. 2016).

massive stars. This is indicated from the periodic X-ray
variation with 5.54 year period. Such variation is widely
observed in wind colliding binaries, with the ellipticalorbit and the enormous stellar winds from both stellar
components. The X-ray intensities become maximum at
the periastron.
At the last periastron of Eta Carinae in 2014 June,MAXI caught the highest X-ray ﬂare among those everrecorded (Negoro et al. 2014). The long X-ray monitoring with MAXI toward the next periastron is highly
deserved.
3.

Summary• 107 ﬂares from 27 low-mass stars were detected.
• dMe stars within 10 pc were all detected.
• No ﬂare has been detected from G type main sequence stars within 20 pc, so far.
• A universal correlation was obtained in the durationof ﬂares vs. the bolometric X-ray luminosity.
• Serendipitously, huge ﬂares were detected from one
un-classiﬁed stellar source.
• The highest X-ray ﬂare ever recorded was detectedfrom Eta Carinae.

References
Maehara, H., Shibayama, T., Notsu, S., et al. 2012,
Nature, 485, 478
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179

Transient Magnetars and Magnetic Field Evolution of Neutron Stars
Teruaki Enoto
The Hakubi Center for Advanced Research and Department of Astronomy,
Kyoto University, Kyoto 606-8302, Japan
E-mail(FS): enoto@kusastro.kyoto-u.ac.jp
Abstract
Accumulating observations of transient activities of magnetized neutron stars have greatly expanded
our knowledge of the magnetic ﬁeld evolution over the neutron star lifetime. In addition to the ordinarytransient magnetars characterized by short bursts and large X-ray luminosity variations, low-ﬁeld magnetars, high-B pulsars, and a compact central object have recently exhibited the similar bursting activity.
Despite their superﬁcial observational diﬀerence, the diversity of neutron stars is expected to be coherently
explained by the evolution of the magnetic ﬁeld as an overriding parameter. In order to observationallyinvestigate the broadband X-ray spectral evolution of the magnetar class, comprehensive re-analyses weremade of all the magnetars observed with Suzaku satellite from 2005 to 2013, combining with early NuSTAR
results and with Swift / RXTE monitorings (Enoto et al. 2017). More than nine sources show the hard Xray power-law radiation above 10 keV, in addition to the well-studied soft thermal component, originatingfrom a stellar surface or vicinity. The X-ray luminosity Lh of the hard component, relative to that of thesoft component Ls , is conﬁrmed to follow the correlation to the dipole magnetic ﬁeld Bd , which results
further reinforced our previous suggestions proposed in Enoto et al. 2010. The Lsd -Lx diagram of transient
sources also implies connections among magnetars, high-B pulsars, and ordinary rotation-powered pulsars,
through activation of hidden magnetic components (higher multipole or toroidal ﬁeld).
Key words: stars: magnetars — stars: magnetic ﬁeld — stars: neutron — X-rays: stars

1.

The Diversity of Magnetized Neutron Stars

Discoveries of new transient magnetars and related families of magnetized neutron stars have greatly expanded
our knowledge of magnetic activities and ﬁeld decay overthe neutron star lifetime. Figure 1 is the up-to-date
P -Ṗ diagram (rotation periods P and their derivativeṖ ) of non-accreting isolated neutron stars. This plot
clearly shows the diversity, called the “neutron star zoo”
(Kaspi 2010; Harding 2013). Majority of this sample
is Rotation-Powered Pulsars (RPPs) which electromagnetic radiation is powered by the spin-down luminosity
Lsd ∝ Ṗ /P 3 , and usually bright in the radio band. The
typical magnetic ﬁeld of the RPPs is B ∼ 1012 G derived from their spin periods and derivatives, assuming
the magnetic dipole radiation in the vacuum. Above the
QED critical ﬁeld Bcr ≡ 4.4 × 1013 G, Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs)
are collectively called magnetars with extremely strong
magnetic ﬁelds of B ∼ 1014 –1015 G (Thompson & Duncan 1995; Mereghetti 2008; Mereghetti et al. 2015).
Their X-ray radiation is believed to be powered by themagnetic energy stored in the stellar interior.
Between the ordinary RPPs and magnetars, there are

growing new classes with intermediate magnetic ﬁeld of
B ∼ 1013 –1014 G. High-B Pulsars (HBP) are radio emitting pulsars with periods at a few hundred millisecond toa few seconds (Olausen et al. 2013). Magnetic ﬁelds of
HBPs are stronger than the ordinary RPPs, almost overlapping the magnetar regime in Figure 1. Their X-ray
luminosities are usually below Lsd , and thus they areclassiﬁed into the RPPs, at least in a quiescent state.
X-ray Isolated Neutron Stars (XINSs) are slowly rotating isolated X-ray objects without radio emission, andthought to be near-by sources (Mereghetti 2011). Magnetic ﬁelds of XINSs derived from their pulsations arealso stronger than the canonical RPPs, and the stored
magnetic energy is thought to aﬀect their hot surface
X-ray radiation. Such a strong ﬁeld was recently supported from a detection of optical polarization fromXINS RX J1856.5−3754 (Mignani et al. 2017).
As shown in Figure 1, the magnetic ﬁeld strengthof neutron star spans more than seven orders of magnitudes, and diﬀerent neutron star families show variety of electromagnetic radiation spectra and timing behaviours. However, these distinct manifestations became
recently blurred when similar magnetar-like activities

180

Bi=1014 G

Bi=1015 G

Bi=1013 G
High-B pulsars withmagnetar outbursts
Bi=1012 G

Low field magnetars
RX J1856.5-3754

Fig. 1. The P -Ṗ diagram of various neutron star families as of 2017 January, based on the catalogues or tables (Manchester et al. 2005; Olausen& Kaspi 2014; Mereghetti 2011), including ordinary pulsars (small circles), magnetar (SGRs and AXPs, large red ﬁlled circles), X-ray isolated
neutron stars (XINSs, orange pendagons), high-B pulsars (HBPs, purple diamonds), central compact objects (CCOs, blue squares), androtating radio transients (RRATs, lightblue triangles). If a pulsar is associated with a supernova remnant (SNRs) or found in a binary
system, they are surrounded by green open circles and blue squares, respectively. The constant spin-down magnetic ﬁelds Bd , characteristic
ages τc , and spin-down luminosity Lsd are shown with dashed lines. The QED critical ﬁeld Bcr = 4.4 × 1013 G and the death line arealso shown (Harding & Lai 2006). The evolutionary paths are overlaid when assuming the empirical Bd -ﬁeld decays (see, e.g., Colpi etal. 2000), calculated with some assumptions at diﬀerent initial magnetic ﬁeld of Bi = 1012 , 1013 , 1014 , and 1015 G. Two HBPs withmagnetar-like outbursts, three activated low-ﬁeld magnetars, and XINS RX J1856.5−3754 are indicated.

181

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1048
2259

1034
1834

0418
1822

=

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d

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1547

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Preliminary

Lc

1030

Lx

2. Transient magnetars and X-ray spectral studies
Some of magnetars suddenly exhibit a large X-ray enhancement, called “outbursts”. In an early phase of Xray outbursts, sporadic short bursts are radiated with
X-ray luminosities exceeding the canonical Eddington
limit (Nakagawa et al. 2007; Enoto et al. 2009; Enoto etal. 2012). These transient magnetic activities, detected
mainly from the Swift satellite, increased the numberof magnetars (see a review, for example, Rea & Esposito2011). The power-law distribution of short bursts ﬂuenceis suggested to show statistical similarity to the solar
ﬂare and seismology on the Earth (Gutenberg-Richter,e.g., Nakagawa et al. 2007). Based on this statisticalsimilarity, the persistent emission is proposed to be composed of accumulation of un-resolved micro short bursts(Nakagawa et al. 2009, but see also, Enoto et al. 2012).
Figure 2 summarizes the observed X-ray luminosities
Lx of the persistent emission, compared with their spindown luminosity Lsd . In the canonical understanding,the persistently bright magnetar Lx exceeds Lsd , while
Lx of the ordinary RPPs stays below Lsd . However,transient magnetars exhibit intermediate property between these two classes. One example is a fast rotating(P ∼2 s) faint AXP 1E 1547.0−5408 which was observedwith Suzaku within a week after the onset of the 2009
outburst (Enoto et al. 2010), and another follow-up observation was further performed one year later (Iwahashi
et al. 2013) with regular Swift monitorings. At an earlyphase, Lx becomes comparable to those of persistently
bright sources, whereas Lx is decreasing back to a quiescence, close to the region of ordinary RPPs (Figure 2).
As shown in this source, some of HBPs or XINSs aredormant magnetars or magnetar descendent.
This view is further reinforced by discoveries of transient magnetar-like activities of three low-ﬁeld magnetars (LFMs) and two HBPs. So far, there are atleast three LFMs known; SGR 0418+5729 (Rea etal. 2013), Swift J1822.3−1606 (Rea et al. 2012), and
3XMM J185246.6+003317 (Rea et al. 2014). Their Bd
strength is comparable to ordinary RPPs (e.g., Bd ∼6 × 1012 G for SGR 0418+5729), while the magnetarlike outburst and short burst activities have been reported. The detection of absorption feature from activated SGR 0418+5729 in soft X-rays, interpreted as

1040

X-ray Luminosity Lx (erg s-1)

have been reported from diﬀerent classes (§2.). Furthermore, magnetic ﬁeld decay and evolution are now
believed to be an overriding key parameter, to unify
the neutron star zoo (Colpi et al. 2000; Viganò et al.
2013). The magnetic ﬁeld decay is also supported froman overestimation of the magnetar characteristic age,when comparing between the pulsar characteristic age
and supernovae remnant age in the X-ray plasma diagnostics. (Nakano et al. 2015).

1028
1028

1030

1032

1034

1036

1038

1040

Spin-down Luminosity Lsd (erg s-1)
Fig. 2. The Lsd -Lx diagram, modiﬁed from Enoto et al. 2017 (seealso Shibata et al. 2016). The ordinary RPPs (brown square symbols) can be powered by the rotational energy (Lx < Lsd ), while
persistently bright magnetars (star symbols) stay in the Lx > Lsd
regime. Transient magnetars (circle symbols or green squares) are
intermediate objects between the above two classes, since Lx ishighly variable, and sometimes crossing Lsd during a decaying
phase of X-ray outbursts.

proton cyclotron resonance, indicates the surface magnetic ﬁeld is much stronger than Bd (Tiengo et al.
2013). We also searched absorption feature for somepersistently bright magnetars with Suzaku, but we only
put upper-limits on the equivalent width of absorption(Miyazaki et al. 2016). As another example, two HBP,PSR J1846−0258 and PSR J1119−6127, also exhibited
transient magnetar activities in 2006 and 2016, respectively (Archibald et al. 2016). Such a transient activity is expected to originated from hidden magnetic ﬁeld
components, for example, higher multipoles on a stellar
surface, or the toroidal ﬁeld in a stellar interior. Usually, only the surface dipole magnetic ﬁeld Bd is measurable from P and Ṗ , while the internal toroidal ﬁeld wasthought to be diﬃcult to be observationally identiﬁed.
However, signatures for the free precession were reported
from two magnetars 4U 0142+61 and 1E 1547.0−5408(Makishima et al. 2014; Makishima et al. 2016). These
are interpreted as the deformation of a neutron star due
to the internal strong toroidal ﬁeld.
All these magnetic activities suggest the magnetic ﬁeld
decay is a key to unify the diﬀerent classes of highly
mangetized neutron stars. Since the Suzaku satellite(Mitsuda et al. 2007), operated from 2005 until 2015,
had the broad-band coverage (0.2–600 keV) using the X-

182

ray Imaging Spectrometer (XIS, Koyama et al. 2007)
and the Hard X-ray Detector (HXD, Takahashi et al.
2007; Kokubun et al. 2007), the soft-thermal surface radiation and hard X-ray magnetospheric emission (Kuiper
et al. 2006) of magnetars were simultaneously observed.
We previously reported that the X-ray luminosity Lh of
the hard component, relative to that of the soft component Ls , follows the positive correlation to the dipole
magnetic ﬁeld Bd (Enoto et al. 2010). Using all themagnetars observed with Suzaku satellite from 2005 to
2013, we performed comprehensive re-analyses of them,
combining with early NuSTAR results and with Swift /RXTE monitorings (Enoto et al. 2017), and re-conﬁrmed
the previous result. This indicates that the magnetar broadband spectra is also governed primarily by the
dipole magnetic ﬁeld intensity Bd .
One of remaining big mysteries is a mechanism
to make various types of highly mangetized neutronstar. Recently, a Compact Central Object (CCO)
1E 161348−5055 at the centre of the supernova remnant RCW 103 showed the magnetar-like activity. In
addition, another mechanism to make a magnetized neutron star, Accretion-Induced Collapse (AIC), is suggested from X-ray observations of symbiotic X-ray binaries (Enoto et al. 2014). These remaining questions
are expected to be answered by future X-ray missions.
In early 2017, Neutron star Interior Composition ExploreR (NICER) is planned to be launched by Space X
Falcon 9 and to be attached to the International Space
Station (ISS) (Arzoumanian et al. 2014). The primary
goal is to determine the equation of state of high density
matter inside neutron stars via precise measurement of
neutron star radii, whereas magnetars and highly mangetized neutron stars are also one of the main targets, especially for transient sources.

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Acknowledge
I gratefully acknowledge S. Takeshige for his supportingof compiling catalogues and references to make ﬁguresin this proceeding. I also thank my collaborators for our
discussion of results presented here, K. Makishima, S.
Shibata, T. Kitaguchi, Y. Suwa, T. Uchide, S. Kisaka,
T. Nakano, H. Murakami, Y. Nakagawa, K. Gendreau,
Z. Arzoumanian, T. Okajima, T. Mihara, W. Iwakiri,
and T. Tamagawa. T. E. was supported by Kakenhi
15H00845 and 16H02198.

Nakagawa, Y. E., et al. 2007, PASJ, 59, 653
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185

Suzaku/WAM Hard X-ray Observations of Solar Flares– from 2005 to 2014
Seiya Yabe,1 M. S. Tashiro,1 Y. Terada,1 T. Yasuda,1 K.Yamaoka,2K.Makishima,3 W. Iwakiri,3 M. Yamauchi,4 N. Ohmori,4
S. Sugita,5 Y. Fukazawa,6 M. Ohno,6 T. Kawano,6 T. Takahashi,7
M. Kokubun,7 Y. Nakagawa,8 Y. Urata,9 K. Nakazawa,10
on behalf of the Suzaku/WAM team.
1

Department of Physics, Saitama University, 255 Shimo-Okubo, Sakura-ku, Saitama-shi,
Saitama 338-8570, Japan
2
Institute for Space-Earth Environmental Research (ISEE), Nagoya University, Furo-cho,
Chikuka-ku, Nagoya, Aichi 464-8601, Japan
3
MAXI Team, The Institute of Physics and Chemical Research (RIKEN), 2-1 Hirosawa,
Wako, Saitama 351-0198, Japan
4
Department of Applied Physics, University of Miyazaki, 1-1 Gakuen kibanadai-nishi, Miyazaki-shi,
Miyazaki 889-2192, Japan
5
Department of Physics, Tokyo Institute for Technology, 2-12-1 Ookayama,
Meguro-ku, Tokyo 152-8550, Japan
6
Department of Physics, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima,
Hiroshima 739-8526, Japan
7
Department of High Energy Astrophysics, Institute of Space and Astronomical Science (ISAS),
Japan Aerospace Exploration Agency (JAXA), 3-1-1 Yoshinodai, Chuo, Sagamihara 252-5210, Japan
8
Advanced Visualization and Computation Research Group, Advanced Earth Information
Research Department, Center for Earth Information Science and Technology (CEIST),
Japan Agency for Marine-Earth Science and Technology (JAMSTEC), Yokohama Institute9
Institute of Astronomy, National Central University, Chung-Li 32054, Taiwan
10Department of Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
E-mail(SY): yabe@heal.phy.saitama-u.ac.jp
Abstract
Solar ﬂare is the largest explosive phenomenon in the solar system, which suddenly releases magnetic
energy built up in the solar atmosphere through reconnection. The rate of solar ﬂare is correlated with thesolar activity whose 11-year period is well known. Solar ﬂare is observed in entire band that we observe.
In particular, hard X-ray radiation is produced by non-thermal electrons accelerated at the magnetic
ﬁeld reconnection, although the detailed mechanism of particle acceleration is still in debate. The Suzaku
Wideband All-sky Monitor (WAM) had observed over 700 solar ﬂares in 50 keV to over MeV band throughthe life of the hard X-ray detector (HXD) from 2005 to 2015. The ﬁrst catalog was published by Endoet al. (2010). They showed the hard X-ray properties of ﬂares in the solar minimum from 2005 to 2009.
Following their study, we carried out systematic spectral analysis of solar ﬂares observed with WAM during
the solar maximum from 2010 to 2014. We recognized no signiﬁcant diﬀerence between the solar maximum
and minimum in the hard X-ray spectrum, and found following ”common” characteristics. First, the hard
X-ray ﬂux of each event is well correlated with the thermal soft X-ray ﬂux. Second, the spectral slopeof the non-thermal hard X-ray component shows clear correlation neither with event duration nor hard
X-ray ﬂux of each event.
Key words: Solar ﬂare: Suzaku/WAM: Hard X-ray:

186

Solar ﬂares are triggered by magnetic ﬁeld reconnection above the photosphere of the sun (Shibata & Magara2011). The magnetically accelerated particles produce
hard X-ray radiation via non-thermal bremsstrahlung(Dennis 1985) and soft X-rays from the heated plasmas
in the magnetic loop follows thermal bremsstrahlung.
Therefore, the hard X-rays are important to know triggering process, but the detailed mechanism of particle
acceleration is still in debate.
Endo et al. (2010) having showed the hard X-ray properties of ﬂares in the solar minimum from 2005 to 2009,
following, we show hard X-ray properties of the solar
ﬂares observed with Suzaku/WAM throughout the 10
years of Suzaku. We use the observation data of the
Wide band All sky Monitor (Yamaoka et al. 2009) thatis active shield of Suzaku (Mitsuda et al. 2007) onboard
Hard X-ray Detector (Takahashi et al. 2007). WAM had
observed 586 solar ﬂares from 2005 to 2014 (Fig.1).

Fig. 2. Correlation plot between averaged hard X-ray ﬂux, duration
time T 90 and soft X-ray ﬂux. Red plots are events in solar maximum, blue plots are events in solar minimum.

Fig. 3. Left : Scatter plot between Hard X-ray ﬂux, duration time
T90 and photon index from WAM spectral. Each data property
are the same as Fig. 2

Fig. 1. Left : Detection number of solar ﬂare by WAM transition.
c NAOJ). The rate of solar ﬂare
Right : Sunspot transition (is correlated with the solar activity. So, we deﬁne from 2006 to2010 as the solar minimum, from 2011 to 2014,2005 as the solar
maximum. Hereafter we divide the observation time in two.

We carried out systematic spectral analysis of 586 events
observed with the WAM from 2005 to 2014.The event criteria are :(1) Simultaneous observation at GOES satellite1 , (2) Detected about 200 keV or more, (3) No Earthoccultation and SAA during the solar ﬂare event. We
analyze thus reduced 274 events (X class 4.4%, M class
48.2%, C class 43.8 %, B class 3.6 %). Observed spectra
are evaluated with a power law function,
A(E) = KE −α ,(1)where K and α are the normalization factor and thephoton index, respecting. We derived T90 2 , 100–300 keV
ﬂux and Photon Index from averaged spectrum of each
event. The derived values are summarized in Fig.2,Fig.3.
The results that we have obtained are summarized as follows. (1) The hard X-ray ﬂux and duration of each event*1 American satellite observing the soft X-ray emission of thesolar ﬂare.
*2 The time duration from the sum of 50–110 keV counts reaches
5 % to 95 % the entire event.

is well correlated with the thermal soft X-ray ﬂux.This isconsistent with the hypothesis that acceleration particles
heat the thermal plasma. (2) The spectral slope of the
non-thermal hard X-ray component shows no signiﬁcantcorrelation with duration,hard X-ray ﬂux. This supports
the idea that long duration ﬂares consist of small scale
ﬂares. (3) We recognize no signiﬁcant diﬀerence in spectral parameters between those in the solar maximum andminimum.
References
Endo, A.,Minoshima, T., Morigami, K., et al. 2010,
PASJ, 62, 1341
Dennis, B. R. 1985, Solphys, 100, 465
Mitsuda, K., Bautz, M., Inoue, H., et al. 2007, PASJ,59, 1
Shibata, K., & Magara, T. 2011, Living Reviews in Solar
Physics, 8, 6
Takahashi, T., Abe, K., Endo, M., et al. 2007, PASJ, 59,
35
Yamaoka, K., Endo, A.,Enoto, T., et al. 2009,PASJ, 61,
35

187

The X-ray luminosities in quiescent phase for superﬂare stars
Ken Yabuki1 , Yohko Tsuboi1 , Ryo Sasaki1 ,Yasuharu Sugawara2 , Satoru Katsuda1 , and Hiroyuki Maehara3
1

Department of Physics, Facility of Science and Engineering, Chuo University,
1-13-27 Kasuga, Bunkyo-ku, Tokyo 112-8551
2
ISAS, JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 229-8510, Japan
3
Okayama Astrophysical Observatory, National Astronomical Observatory of Japan,
3037-5 Honjo, Kamogata, Asakuchi, Okayama, Japan, 719-0232
E-mail(KY): yabuki@phys.chuo-u.ac.jp
Abstract
Maehara et al. (2012) and Shibayama et al. (2013) have searched for superﬂares on G-type dwarfs using
Kepler data in the optical band. They found that slowly rotating stars, with a spin period of about 20 days
like the Sun, caused superﬂares. It is believed that slowly rotating stars have a low X-ray luminosity anda small magnetic ﬂux. Hence the superﬂares that occur on slowly rotating stars are strange phenomena.
Here, we have searched for Lxq of superﬂare stars, using available X-ray archival data. As the result, wefound that 11 stars are in the ﬁeld of view of XMM-Newton and 7 of them are detected. In addition,
we conﬁrmed that 3 stars are identiﬁed with sources in the ROSAT All-Sky Survey Faint Source Catalog(Voges et al. 2000). We discovered a positive correlation between Lxq and the maximum ﬂare energy
during the Kepler observation (1.4 year) Etot,opt : Etot,opt ∝ L1.2
xq . In addition, there is no correlation
between Lxq with Prot for superﬂare stars. These results indicate that the ﬂare mechanism is common
while the generation mechanism of the magnetic ﬁelds is diﬀerent.
Key words: stars:ﬂares—stars:G type dwarfs

1.

Introduction

The X-ray emission from stars is associated with the stellar magnetic ﬁeld. In fact, Pevtsov et al (2003) reported
that the quiescent X-ray luminosity are proportional tothe magnetic ﬂux. In addition, the quiescent X-ray luminosity declines with increasing rotation periods (Prot )for the stars with Prot of more than 7 days. Hence slowly
rotating stars have a small magnetic ﬂux. Since ﬂares arecaused by releasing magnetic energy, it is believed thatslowly rotating stars cannot cause large ﬂares.
However, the large ﬂares occurring on slowly rotatingstars were reported by Maehara et al. (2012). Maehara
et al. (2012) searched for superﬂares on G-type dwarfs in
optical white light using Kepler data and found superﬂares on slowly rotating stars with a spin period of about
20 days like the Sun. They derived the total bolometric energy of the superﬂares and found that the largest
values of those for each object do not show any clear
correlations with the period of the stellar rotation. This
result indicates a diﬀerent trend from that obtained in
X-rays. Therefore we investigate the correlation between
the maximum energy of white light ﬂares during Kepler
observation (1.4 years)(Etot,opt ) and quiescent X-ray lu-

minosity (Lxq ) for each object.
2. Analysis and result
Target stars are 279 superﬂare stars reported in
Shibayama et al. (2013). To ﬁnd superﬂare stars detected in X-rays, we searched for X-ray archival data.
As a result, 11 stars were located in the ﬁeld of view of
XMM-Newton and 7 stars were detected. 3 stars wereidentiﬁed with the sources in ROSAT All Sky Survey
faint Source catalog. One of the three stars was also
detected with XMM-Newton. Except for these stars, 4stars were in the ﬁeld of view of Chandra, although nostars were detected. In total, we found 9 superﬂare stars
detected in X-rays. We investigate these stars in thiswork.
First, we made a light curve for each star to study
time variations. We ﬁt each light curve with a constant.
As a result, the reduced χ2 is less than 2.0 for each star.
Therefore we consider that these stars are in the quiescent phase. Second, we calculated Lxq . For the stars
detected with signiﬁcance levels larger than 7σ, we ﬁt
the spectra with a thin thermal plasma model to calculate Lxq . For the other stars, we calculated Lxq by using

28
29
30
log(Quiescent X−ray Luminosity [erg/s])

31

Fig. 1. The relation between Etot,opt and Lxq . The vertical axis
is the maximum total ﬂare energy in the observation period (1.4
years) and the horizontal axis is Lxq .

webpimms assuming a thin thermal plasma model; we assume NH , kT and the abundance as 1018 cm−2 , 1.0 keV,
and 0.4, respectively. Consequently, Lxq distributed at
the range of 1029 –1031 erg s−1 . Figure 1 shows the relation between Etot,opt and Lxq . We found the following
relation:
1.2+0.3

Etot,opt ∝ Lxq −0.4

(1)This equation implies that the maximum energy of whitelight ﬂares is determined by the quiescent X-ray luminosity.
3. Discussion
3.1. Comparing with X-ray relation
In X-rays, Sasaki et al. (2017) investigated the relation
between the maximum energy of the X-ray ﬂares (Etot,x )and Lxq using MAXI data. They found the following
relation:
Etot,x ∝ L1.26
xq

±0.04

(2)

In spite of the diﬀerent energy ranges, both Equation(1) and (2) are roughly proportional to Lxq , and the index of Lxq are the same within the error ranges. Thismeans that the ratio of the maximum ﬂare energy between white light and X-rays is constant even though
the quiescent X-ray luminosity is diﬀerent.
The target stars of Sasaki et al. (2017) are several
types of stars (RS CVn systems, Algol systems, dMe,
dKe, Young Stellar Object and K-type variable star). On
the other hand, our target stars are only G-type dwarfs.
Although target stars are diﬀerent, we obtained a similar

29
28
27

Log(Xíray Luminocity[erg/s])

36
35
34

log(Total flare energy [erg])

30

37

188

1

10
Period[day]Fig. 2. The relation between Lxq and the rotational period. The
vertical axis is Lxq and the horizontal axis is the rotational period.
The stars and the circles indicate superﬂare stars and the G-typestars on which superﬂares have not been detected, respectively.

relation with Sasaki et al. (2017). This implies that
Equation (1) holds for several types of stars.
4. The diﬀerence of superﬂare stars and non-superﬂarestars
It is widely known that Lxq correlates with Prot among
G-type stars on which superﬂares have not been detected(we called these stars as non-superﬂare stars) (Pizzolato
et al. 2003). On the other hand, among superﬂare stars,
Lxq do not correlate with Prot (Figure 2).
In general, rapidly rotating stars generate strong magnetic ﬁelds. If the generation mechanism for the magnetic ﬁelds of superﬂare stars is the same as nonsuperﬂare stars, we can predict that superﬂare stars
have the same Lxq –Prot relation as non-superﬂare stars.
However, our result is inconsistent with this prediction.
Hence, we can consider that Figure 2 indicates the possibility that the generation mechanism for the magnetic
ﬁelds of superﬂare stars is diﬀerent from non-superﬂare
stars.
References
Maehara, H., et al. 2012, Nature, 485, 478
Pevtsov, A., et al. 2003, ApJ, 598, 1387
Pizzolato, N., et al. 2003, A&A, 397, 147
Sasaki, R., et al. 2017, in these proceedings
Shibayama, T., et al. 2013, ApJS, 209, 5

189

Statistical research of hyper X-ray ﬂares detected with MAXI
– a correlation between the quiescent luminosity and the largest ﬂare energy –Ryo Sasaki,1 Yohko Tsuboi,1 Satoru Katsuda,1 Ken Yabuki1
Yumiko Nakamura1 Yasuharu Sugawara2 Masaru Matsuoka3 and the MAXI team1

Department of Physics, Facilty of Science and Engineeriong, Chuo University,
1-13-27 Kasuga, Bunkyo-ku, Tokyo 112-8551, Japan
2
ISAS, JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 229-8510, Japan
3
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(RS): sasaki@phys.chuo-u.ac.jp
Abstract
MAXI started its operation in 2009 August. Thanks to its unprecedentedly high sensitivity as an allsky X-ray monitor and to its capability of real-time data transfer, we have detected 106 strong ﬂares fromtwenty-seven active stars (fourteen RS CVn systems, one Algol systems, nine dMe stars, one dKe star, one
Young Stellar Object and one K-type variable star). X-ray energies of these ﬂares were 6 × 1033 –9 × 1038
erg in the 2–20 keV band. These ﬂares can be high ends of their own categories (Tsuboi et al. 2016).
Quiescent state X-ray emissions is thought to originate from the magnetic activity as well as stellar ﬂares.
In response to this, we compared the largest X-ray ﬂare energy in 7.25 years (running period of MAXI
until 2016 November) for each object (Etot,max ) with quiescent state X-ray luminosity (Lx,q ). Each Lx,q
was extracted from ROSAT All-Sky Survey Bright Source Catalogue (Voges et al.1999). As a result, wediscovered a correlation of Etot,max ∝ L1.26
x,q , for the ﬁrst time. With this correlation, we can predict thehighest ﬂare energy in 7.25 years from Lx,q for all stars.
Key words: stars:ﬂares — stars:RS-CVn type — stars:dMe

1. Introduction
Stellar ﬂares are caused by magnetic reconnection
which is speculated by numerous preceding solar studies. Flares often accelerate electrons and protons over
MeV and GeV, respectively. The particles precipitate
along the magnetic ﬁelds into the chromosphere, suddenly heating the plasma at the bottom of the magnetic
loop up to very high temperatures.
A large number of solar and stellar ﬂares were detected
so far. MAXI detected many stellar ﬂares. These ﬂares
are located at the high ends in correlation for both theduration time (τ ) vs. X-ray luminosity (Lx ) and theemission measure vs. temperature (Tsuboi et al. 2016).
Preceding works have discovered that a distribution ofstellar ﬂare energy can be approximated by a power-low
Nτobs(E > Etot ) = A E−α
tot

year−1 ,(1)

where N(E > Etot ) and τobs are the number of eventswith energy larger than the ﬂare energy (Etot ) and observed years, respectively. The index α is 0.6 on RS

CVn system (Osten and Brown 1999), 0.52 and 0.7 on
M dwarfs (Collura et al. 1988 and Pallavicini et al. 1990).
Quiescent state X-ray emissions is thought to originate
from the magnetic activity as well as stellar ﬂares. In
response to this, the distribution of stellar ﬂare, which
ought to be typiﬁed by the largest ﬂare energy in a certain period of time, can be considered to be related withquiescent X-ray emissions. To resolve the relevance, we
focused on the largest ﬂare energy which detected withMAXI in the running period of MAXI.
2. Result
MAXI has been observing the all sky in 7.25 years whichis from 2009 August to 2016 November. MAXI has detected 106 strong ﬂares from twenty-seven active stars(fourteen RS CVn systems, one Algol system, nine dMestars, one dKe star, one Young Stellar Object and one
K-type variable star). The locations of these stars on
all-sky map are shown in Figure 1. The energies of theseﬂares range from 6 × 1033 to 9 × 1038 erg in the 2–20 keVband, as shown in Figure 2. The ﬂare energies detectedwith MAXI are more than two orders of magnitude larger

40

190





36



 
 




















32

Fig. 1. Location of ﬂare stars detected with MAXI on all-sky map
28







34

log (Etot,max [erg])

38




 

30
log (Lx,q [erg s−1])

 

 
 
 
 



32

Fig. 3. The correlation between Etot,max and Lx,q .

Fig. 2. The ﬂare energy histgram detected with MAXI

than the maximum solar ﬂare (≈ 1032 erg).
Fig. 4. The diﬀerence of the ﬂare energy distribution of each star.
Star A, B, C are arranged in order from a smaller Etot,max .

3. Discussion
3.1. The maximum ﬂare energy prediction
The obtained ﬂares can be high ends of their own categories (Tsuboi et al. 2016). We compared the maximum
ﬂare energy in running period of MAXI for each object (Etot,max ) with quiescent X-ray luminosities (Lx,q ).
Flare energies are derived by multiplying Lx by τ . The
value of τ were not measured for seven stars. For theseﬂares, we used a correlation of τ ∝ L0.2
x that was discovered in Tsuboi et al. (2016). We plotted the seven stars
by open mark. On the other hand, each Lx,q was extracted from ROSAT All-Sky Bright Source Catalogue(Voges et al. 1999). We ﬁtted all the data in Figure 3 bya power-law model and obtained the best-ﬁt function ofEtot,max = 10−2.04±1.12 L1.26±0.04
x,q

erg.

(2)

The error of the coeﬃcient and the power are 1–σ conﬁdence level. The best-ﬁt model is shown in Figure 3.
With this correlation, we are able to predict the maximum X-ray ﬂare energy of individual stars during theobserving span of 7.25 years from their Lx,q .
3.2. Flare energy distribution
The equation 2 implies that ﬂare mechanisms among theall stellar systems are intrinsically the same. Assumingthat all stellar ﬂare systems obey the same index “α”

of equation 1, we ﬁxed the index to 0.6. The N is 1,
because Etot,max occurred once in τobs (7.25 years). The
schematic view of this idea is shown in Figure 4. The
equation of the number of ﬂares is
1(E ≥ Etot,max ) = A E−0.6
tot,max
7.25 years

year−1 . (3)Etot,max With the equation 2, we were able to derive thenormalization “A”.
A = 1.38 × 1019.38 × (

Lx,q
)0.72
1030 erg s−1(4)

In this way, we also succeeded in predicting the number of ﬂares in 7.25 years of each star from the quiescentstate X-ray luminosity.
References
Collura, A. et al. 1988, A&A, 205, 197
Osten, R A., Brown, A. 1999, ApJ, 515, 746
Tsuboi, Y. et al. 2016, PASJ, 68, 90
Voges, W. et al. 1999, A&A, 349, 389

191

Super ﬂare candidates detected by XMM-Newton
Y. Nakamura,1 Y.Tsuboi,1 Y.Terashima,2
Y.Sugawara3 and S.Katsuda1
1

Department of Physics, Facilty of Science and Engineering, Chuo University,
1-13-27 Kasuga, Bunkyo-ku, Tokyo 112-8551Chuo University
2
Department of Physics, Ehime University, Matsuyama, Ehime 790-8577, Japan
3
ISAS, JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 252-8510, Japan
E-mail(YN): y.nakammura@phy.chuo-u.ac.jp
Abstract
While superﬂare surveys on main-sequence stars have been well performed with Kepler satellite in
the optical band (e.g. Maehara et al. 2012), there are few reports in the X-ray band. Hence necessityof statistical studies in the X-ray band has been enhanced recently. We found 23 objects which havethe light curves like those of stellar ﬂares in 2XMMi-DR3 catalog (Watson et al. 2009). Among them,
we have identiﬁed 22 objects with the stars listed in available optical and infrared catalogs. We madeSEDs from optical and infrared data, and derived temperatures and spectral types (M-type: 18, K-type: 3,
F-type: 1). Temperatures of nine objects are lower than 3000 K. These nine stars are candidates of browndwarf. To determine distance and age, we compared proper motion of our objects with those of near–earth moving group members. One M-type object is concluded to be a new member of AB Dor moving
group (distance: 7–77 pc, age: 50–120 Myr, Malo et al. 2012), from the coordinate and proper motion.
According to observation of Gaia, distances of two objects (K-type: 1, F-type: 1) are 210 (198–222) pc,
833 (189–1477) pc. Then we estimated absolute magnitude from the distances. Absolute magnitude of
the K-type star is brighter than main-sequence K-type stars, while the F-type star has possibility of beinga main-sequence star. Assuming that other objects (K-type: 2, M-type: 8 [Temperature > 3000 K]) arein the main sequence phase, we calculated their lower limits of distances from apparent magnitude andthe temperature. The lower limits of distances of them are ranging also 10–200 pc. We calculated X-ray
luminosities of thirteen objects from distances. Six of the thirteen objects have luminosities at least 1029
erg s−1 , well in the superﬂare class, suggesting that they are super ﬂare star candidates. Furthermore,
there is a possibility that four objects (F-type: 1, K-type: 2, M-type: 1) of the six superﬂare stars are
main-sequence stars. We were able to detect four super ﬂare star candidates in main-sequence phases.
Key words: Stars: ﬂares
1. Introduction
Flares ten times more energetic than the largest solar
ﬂares are called super ﬂare. Generally, superﬂares havebeen observed on binaries and YSOs in X-rays. Whereas
superﬂare surveys for main-sequence stars have been performed actively in optical (e.g. Maehara et al. 2012), statistical studies in X-rays are still scarce.
2. No bias ﬂare search
Our sample was selected from 2XMMi-DR3 produced by
the XMM Survey Science Centre (Watson et al. 2009). In
order to search for highly variable objects, we performed
chi-square tests with a null hypothesis that the objects
have constant ﬂuxes. These probabilities were calculated
from the time series in 0.2–12 keV by using the Sci-

ence Analysis System (SAS) task ekstest excluding highbackground ﬂaring times. We selected sources satisfying the following conditions: (1) the probability calculated by using EPIC-pn data < 10−5 , (2) the count rate
for EPIC-pn in 0.2–12 keV > 0.03 counts s−1 , and (3)Galactic latitude |b| >10, where all of these parameters
are listed in 2XMMi-DR3. 1100 sources fulﬁlled thesecriteria. We applied further ﬁltering conditions to construct our sample. We discarded sources in the ﬁelds ofstar forming regions, Small Magellanic Clouds, or Large
Magellanic Clouds. Sources with object types Galaxy,
X-ray source, or unidentiﬁed shown in the NASA/IPAC
Extragalactic Database (NED) or the SIMBAD were regarded as AGN candidates. We selected stars which havethe light curves resemble those of stellar ﬂares (Abrupt




















   

  



192









Fig. 1. Distribution of temperatures (M-type:18, K-type:3, F-type:1)

rise of ﬂux followed by exponential decay). After all
screening, 23 sources were ﬁnally selected. Of these, 22
objects are found to have counterparts in optical (USNOB1,GSC2.3.2) and infrared (2MASS, WISE) catalogs.
3. Discussion
3.1. Temperature
We made SEDs from optical and infrared data. Then we
determined temperatures and spectral type (M-type: 18,
K-type: 3, F-type: 1, shown in Fig 1). Temperatures ofnine objects are cooler than 3000 K. We consider theseobjects to be brown dwarf candidates, judging from thetemperatures (Thackrah et al. 1997). It should be notedthat they may be members of a rare class of old browndwarfs that show ﬂares in the X-ray band (e.g. Rutledge
et al. 2000, Age: 500 Myr). Main-sequence stars withintermediate spectral types are considered intrinsically
X-ray inactive. A detailed investigation on the F-type
object is necessary.
3.2. Comparison with near-earth moving group
We compared proper motions of our objects with propermotion of near–earth moving group to determine distances and age. One of our objects has proper motion
resembling those of AB Dor moving group (Distance: 7–
77 pc, Age: 50–120 Myr ; malo et al. 2012). From thisreason, we concluded that the object is a new memberof AB Dor moving group. Proper motion of AB Dor
Moving group and our sample are shown in Fig 2. AB
Dor Moving is in YSO phase. YSOs produce ﬂares frequently (e.g. Imanishi et al. 2003). The young age of AB
Dor moving group also supports the membership.
3.3. DistanceWe determined distance of two objects (K-type: 210[198–222] pc, F-type: 833 [189–1477] pc) based on parallaxes measured with Gaia, and we estimated absolutemagnitude. Absolute magnitude of the K-type star isbrighter than main-sequence K-type stars, while F-typestar has possibility of being a main-sequence star. Assuming that other nine objects (Temperature > 3000 K)
are in the main sequence phase, we determined lower 

        Fig. 2. Proper motions of AB Dor Moving group and our object.
Solid line and dotted line show proper motion of AB Dor moving
group and our object.

Fig. 3. Lower limits of distances of the ten samples.

limits of distances from apparent magnitude and thetemperatures. We used V-band magnitude listed in
GSC2.3.2 catalog as apparent magnitude, while, for objects not listed in that catalog, we estimated V-band
magnitude from SED. We show lower limits of distancesin Fig 3.
3.4. X-ray Luminosity
We calculated X-ray luminosities of our objects fromdistances (A member of AB Dor moving group: 1, With
Gaia: 2, lower limits: 10). Luminosities were shown in
Fig 4. Six ﬂares have X-ray luminosities at least 1029
ergs s−1 , well in super ﬂare class.
In conclusion, we succeeded in detecting six super ﬂarestars. Furthermore, there is a possibility that four objects (F-type: 1, K-type: 2, M-type: 1) of the six superﬂare stars are main-sequence stars. We were able to detect super ﬂare star candidates in main-sequence phases.

   





Fig. 4. X-ray luminosities of our objects.

References
Maehara et al. 2012, Nature, Volume 485
Watoson et al. 2009, A&A, 493, 339
Thackrah et al. 1997, MNRS, 284, 507T
Malo et al. 2013, APJ, 762, 88
Imanishi et al. 2003, PASJ, 55, 653
Rutledge et al. 2000, APJ, 538L, 141R

193

X-ray spectral variability of the colliding wind binary WR140– the origin of the cool plasma component –Yasuharu Sugawara,1 Yoshitomo Maeda,1,2 Kenji Hamaguchi3,4 and Yohko Tsuboi5
1Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency,
3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 229-8510, Japan2
Department of Space and Astronautical Science, SOKENDAI (The Graduate University for
Advanced Studies), 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 252-5210, Japan
3
CRESST and X-ray Astrophysics Laboratory NASA/GSFC, Greenbelt, MD 20771, USA
4
Department of Physics, University of Maryland, Baltimore County,
1000 Hilltop Circle, Baltimore, MD 21250, USA
5
Department of Physics, Faculty of Science & Engineering, Chuo University,
1-13-27 Kasuga, Bunkyo, Tokyo 112-8551, Japan
E-mail(YS): sugawara.yasuharu@jaxa.jp
Abstract
We report the results from the X-ray spectral analysis of the colliding wind binary WR140(WC7pd+O5.5fc) using archival XMM-Newton, Suzaku and Chandra data. Recently, we reported that acool plasma component in a recombining phase was discovered from the WR140 X-ray spectrum at near
periastron. In order to uncover the origin of the cool plasma component, we performed the spectral analysis
using XMM-Newton, Suzaku and Chandra data at near periastron. These observations cover four diﬀerent
epochs from 2008 December 26 to 2009 January 25 for a total exposure of about 184 ks. The spectra of thecool plasma component were well-ﬁtted by a single-absorbed non-equilibrium ionization collisional plasma
model. It is found that the emission measure of the cool plasma component decreased at approximately
ﬁfty percent per month. This variation indicates that the plasma was cooled down by radiation and/or was
expanded. As one interpretation, the cool component may be a relic of the wind-wind collision plasma,
and may represent a transitional phase from the compressed hot gas to dust formation.
Key words: stars: Wolf-Rayet ― binaries: general ― stars: winds, outﬂows ― X-rays: individual(WR140)1.

Introduction

Massive binaries composed by a Wolf-Rayet star and a
OB star have the highest temperature plasma, producedby the collision of the winds. Colliding wind binary is thebest testing ground for plasma shock physics, because
plasma properties vary with binary separations. Long
period binary WR140 (WC7pd+O5.5fc, P=7.94 yr) is
considered as the textbook example of an episodic dustmaking colliding wind binary and a good natural laboratory for the study of shock physics, because the orbitalparameters have well determined.
The cool plasma component was discovered for theﬁrst time by WR140 Suzaku observations around 2009
periastron (Sugawara et al. 2015). The simple analysis of
one-temperature collisional equilibrium plasma emission(APEC: Smith et al. 2001) for the cool component failed
to reproduce the observed spectra with the residuals at

1.21 keV and other energies remaining in the spectral
ﬁtting. On the other hand, the spectrum of the cool
plasma component is reproduced well with the recombining collisional plasma model (vrnei1 in Xspec).
In order to uncover the origin of the cool plasma component, we performed the spectral analysis using XMMNewton, Suzaku and Chandra data at near periastron.
2. Analysis and Results
We used archived XMM-Newton/EPICs, Suzaku/XISs,
and Chandra/ACIS-I datasets around 2009 periastronpassage. The observations of WR140 were conductedwith XMM-Newton on 2008 December 26 (ObsID
0555471001), with Suzaku on 2009 January 04–05 and*1 See http://heasarc.gsfc.nasa.gov/xanadu/xspec/manual/
XSmodelRnei.html

194 
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Fig. 1. The background subtracted spectra of WR140 for each phase and the ﬁtting models. The left, center and bottom panels show the
spectra of XMM-Newton/EPICs, Suzaku/XISs, and Chandra/ACIS-I, respectively. The red and blue lines in the upper panels show thecool and hot components, respectively. The lower panels show the residuals of the data from the model.

3. Discussion
We found for the ﬁrst time that the emission measure
and the absorption column density for the cool component decreased. A potential origin of the cool plasma
is as follows. The plasma was heated in the past by awind-wind collision shock. The heated plasma then escaped from the dense wind region, and is recombining*2 We adopted two absorption components. One is the absorption component (varabs) for the W-R wind. Another is aninterstellar absorption component, with elemental abundances
ﬁxed at ISM abundances (TBabs: Wilms et al. 2000 ).

XMM

SuzakuChandranH [1021 cm−2]

E.M. [1054 cm−3]
2
5
10

13–15 (Sequence number 403032010 and 403033010), andwith Chandra on 2009 January 25 (ObsID 9918), respectively. The total exposure time was about 184 ks. The
Suzaku datasets were summed up for two observations
in order to get the good photon-statistic spectrum.
In all spectra, we found a soft emission component below 2 keV reported as cool component. We ﬁtted thespectra for each phase with a two-temperature plasma
model with independent absorption components (ﬁgure 1). One plasma model is the dominant hot component above 2 keV is derived from the wind-wind collision (model: TBabs*varabs*vAPEC)2 , and the othermodel is the cool component below 2 keV (model:
TBabs*vrnei). As for the hot plasma, the elemental
abundance were ﬁxed at the reported values (table 3 in
Sugawara et al. 2015). On the other hand, as for the cool
plasma, the plasma parameters were ﬁxed at reported
values (table 2 in Sugawara et al. 2015) except for theemission measure and the absorption column density.
Figure 2 shows the variation of the emission measure and the column density for the cool component.
The emission measure decreased at approximately ﬁfty
percent per month. The absorption column density of
Suzaku spectrum was decreased at approximately thirty
percent relative to that for XMM-Newton spectrum.

2.995

3

3.005

phase
Fig. 2. The variation of the emission measure (red circle) and theabsorption column density (blue square) for cool component withorbital phase. Orbital phases were derived using parameters given
in Monnier et al. (2011) The dashed line shows a periastron phase.

electrons and emitting the radiative recombination continuum along with the collisional plasma emission. In
this case, the cool plasma component may represent atransitional phase from the compressed hot gas to dust
formation.
This research has made use of data and/or software provided by the High Energy Astrophysics Science Archive
Research Center (HEASARC), which is a service of the
Astrophysics Science Division at NASA/GSFC and theHigh Energy Astrophysics Division of the Smithsonian
Astrophysical Observatory This work was supported by
JSPS KAKENHI Grant Number JP16K17667.
References
Monnier, J. D., et al. 2011, ApJ, 742, L1
Smith, R. K., et al. 2001, ApJL, 556, L91
Sugawara, Y., et al. 2015, PASJ, 67, 121
Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914

195

superﬂares on Sun-like stars
Kazunari Shibata1
1

Kwasan and Hida Observatories
E-mail(KS): shibata@kwasan.kyoto-u.ac.jp
Abstract
Many stars show ﬂares similar to solar ﬂares, and often such stellar ﬂares are much more energetic thansolar ﬂares. The total energy of a solar ﬂare is typically 1029 –1032 erg. There are much more energetic
ﬂares (1033 –1038 erg) in stars, especially in young stars with rapid rotation. These are called superﬂares.
We propose that these stellar superﬂares can be understood in a uniﬁed way based on the reconnection
mechanism which has been developed to explain solar ﬂares. Recently, it has been revealed that superﬂares
with energy of 1034 –1035 erg (100–1000 times of the largest solar ﬂares) occur with frequency of once in
800–5000 years on Sun-like stars with slow rotation, which are similar to our Sun. These superﬂares areusually associated with large spots with area A = 103 –105 in unit of one millionth of solar hemisphere,
much larger than normal sunspots (with area A = 100–1000) on the Sun. It has become clear thatsuperﬂares can occur on slowly rotating Sun-like stars because very large star spots can be generated inthese slowly rotating stars, though frequency is very small. Hence, the problem of superﬂare occurrence
becomes dynamo problem; why can a very large star spot be generated in slowly rotating stars like our
Sun?

196

197

ActiveGalactic
Nuclei

198

199

Hard X-ray Luminosity Function of Tidal Disruption Events: First
Results from MAXI Extragalactic Survey
Taiki Kawamuro,1 Yoshihiro Ueda,1 Megumi Shidatsu,2 Takafumi Hori,1Nobuyuki Kawai,3 Hitoshi Negoro,4 and Tatehiro Mihara2
1

Department of Astronomy, Kyoto University, Kyoto 606-8502, JapanMAXI team, RIKEN, 2-1, Hirosawa, Wako-shi, Saitama 351-0198, Japan3
Department of Physics, Faculty of Science Tokyo Institute of Technology, Tokyo 152-8551, Japan
4
Department of Physics, Nihon University, 1-8-14 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8308, Japan
E-mail(TK): kawamuro@kusastro.kyoto-u.ac.jp
2

Abstract
We derive the ﬁrst hard X-ray luminosity function (XLF) of stellar tidal disruption events (TDEs) by
supermassive black holes (SMBHs), which gives an occurrence rate of TDEs per unit volume as a functionof peak luminosity and redshift, utilizing an unbiased sample observed with the MAXI . On the basis of
the light curves characterized by power-law decay with an index of −5/3, a systematic search using the
MAXI data in the ﬁrst 37 months detected four TDEs, all of which have been found in the literature. To
formulate the TDE XLF, we consider the mass function of SMBHs, that of disrupted stars, the speciﬁc
TDE rate as a function of SMBH mass, and the fraction of TDEs with relativistic jets. We perform an
unbinned maximum likelihood ﬁt to the MAXI TDE list. The results suggest that the intrinsic fractionof the jet-accompanying events is 0.0007%–34%. We conﬁrm that at z  1.5 the contamination by TDEs
to the hard X-ray luminosity functions of active galactic nuclei is not signiﬁcant and hence that theircontribution to the growth of SMBHs is negligible at the redshifts.
Key words: galaxies: active — X-rays: galaxies

1.

Introduction

Tidal disruption events (TDEs) by supermassive black
holes (SMBHs) occur when a star gets close to a SMBH
enough for the tidal force to exceed the self-gravity ofthe star. A part of the disrupted star subsequently accretes onto the SMBH, which is observed as a luminous
ﬂare. These suggest that the TDE contributes to themass growth of SMBHs. To quantitatively discuss theeﬀect of TDEs on the growth history, luminosity dependence of the TDE rate, or “luminosity function”, has tobe constrained based on a statistically complete sample.
Some studies have theoretically estimated the occurrence rate of TDEs to be 10−5 –10−4 galaxy−1 yr−1 (e.g.,
Magorrian & Tremaine 1999). Many observational results (e.g., Donley et al. 2002) are in rough agreementswith the predicted TDE rate. Furthermore, the dependence of the rate on SMBH mass is numerically calculated (e.g., Stone & Metzger 2014). Considering that theﬂare luminosity of a TDE depends on the SMBH mass, it
is possible theoretically to estimate the luminosity function of TDEs (Milosavljević et al. 2006). However, observational studies that directly constrain the TDE lumi-

nosity function based on a statistically complete sample
have been limited so far.
X-ray surveys covering a large sky area are very useful to detect TDEs, and to make the complete sample,
because we cannot predict when and where an event occurs. In fact, wide-area X-ray surveys performed withROSAT , XMM-Newton, INTEGRAL, and Swift (e.g.,
Komossa & Bade 1999; Burrows et al. 2011; Nikolajuk &Walter 2013) have discovered many of the TDEs. Sometimes accompanying relativistic jets were reported. The
identiﬁcations of TDEs were mainly based on their variability characteristics, such as a large amplitude and theunique decline law of the light curve (e.g., Komossa &Bade 1999).
In this proceeding, we present the derivation of thehard X-ray luminosity function (XLF) of TDEs, using astatistically complete sample. Based on the MAXI mission, we make the sample by systematically searchingfor hard X-ray transient events at high galactic latitudes(|b| > 10◦ ), and identify TDEs. We then derive the XLF
of TDEs associated with and without relativistic jets individually. Eventually, the result enables us to estimate

200

the contribution of TDEs to the mass growth of SMBHs.
Note that more details are presented in Kawamuro et al.
(2016).
Throughout this proceeding, we assume a Λ cold darkmatter model with H0 = 70 km s−1 Mpc−1 , ΩM = 0.3,
and ΩΛ = 0.7. The “log” denotes the base-10 logarithm,
while the “ln”, the natural logarithm.
2. Search of TDE From MAXI Data
2.1. Observations and Data Reduction
In this work, we search for TDEs by utilizing the
MAXI /GSC data (Matsuoka et al. 2009; Mihara et al.
2011). The MAXI /GSC has two instantaneous ﬁelds-ofview of 1◦ .5×160◦ separated by 90 degrees. By rotating
according to the orbital motion of the international spacestation, it covers a large fraction of the sky (95%) in one
day (Sugizaki et al. 2011). We analyze the data takenin the ﬁrst 37 months since the beginning of the operation (from 2009 September 23 to 2012 October 15).
We also restrict our analysis to high galactic latitudes(|b| > 10◦ ). Hence, the data we use are the same as
those made for creating the second MAXI /GSC catalog(Hiroi et al. 2013), which contains 500 sources detectedin the 4–10 keV band from the data integrated over thewhole period. The details of the data selection criteria
are described in Section 2 of Hiroi et al. (2013).
2.2. Identiﬁcation of TDEs in MAXI Catalogs
To detect TDEs whose typical time scale of months toyears (e.g., Komossa & Bade 1999) as completely as
possible, we newly construct the MAXI /GSC “transient
source catalog”. We optimize the analysis to ﬁnd variable objects on the time scale of 30 or 90 days. Namely,
we split the whole data into 30 or 90 day bins, and independently perform source detection in each dataset tosearch for new sources that are not listed in the 2nd
MAXI catalog (Hiroi et al. 2013). The detail can befound in Appendix 1 of Kawamuro et al. (2016). As aresult, we detect 10 transient sources with the detection
signiﬁcance sD > 5.5 in either of the time-sliced datasets,
where sD is deﬁned as (best-ﬁt ﬂux in 4–10 keV)/(its 1σstatistical error). The sensitivity limit for the peak ﬂux
averaged for 30 days is ∼ 2.5 mCrab.
On the basis of positional coincidence, we identify 3
TDE candidates in the MAXI transient source catalog
from the literature, Swift J1112.2–8238 (hereafter SwiftJ1112–82; Brown et al. 2015), Swift J1644+57 (Burrows
et al. 2011), and Swift J2058.4+0516 (hereafter Swift
J2058+05; Cenko et al. 2012). We also identify another candidate that occurred in NGC 4845 (Nikolajuk& Walter 2013), which has been already listed in the 2nd
MAXI /GSC catalog (Hiroi et al. 2013). The detailed information of each TDE from the literature is summarizedin Table 1.

Table 1. Our Sample of Tidal Disruption Events
Name[1]
Swift J1112.2–8238
Swift J164449.3+573451
Swift J2058.4+0516
NGC 4845

z
[2]
0.89
0.354
1.18530.004110log LX
[3]
47.1
46.6
47.5
42.3

Γ
[4]
10
>2
-

δ
[5]
16
-

M∗
[6]
0.15
0.1
0.02

Notes.
Col. [1]: Name of the TDE. Col. [2]: Redshift. Col. [3]: Luminosity (erg s−1 ) in the 4–10 keV band. Col. [4]: The Bulk Lorentz
factor. Col. [5]: The Doppler factor. Col. [6]: Accreted mass inunits of solar mass.

2.3. Search for Unidentiﬁed TDEs
To derive the XLF, it is very important to perform acomplete survey at a given ﬂux limit. Hence, we searchfor other possible TDEs that are not reported in theliterature from these MAXI catalogs. We make use of
the general characteristics of the light curve pattern of
TDEs: rapid ﬂux increase and power-law decay with anindex of –5/3 as a function of time (Phinney 1989). For
this purpose, we make the light curves of all sources in
the 2nd MAXI /GSC catalog (Hiroi et al. 2013) and in
the transient catalog in 10 days, 30 days, and 90 days
bins, in three energy bands, 3–4 keV, 4–10 keV, and 3–
10 keV. The ﬂuxes in each time bin are obtained by thesame image ﬁtting method as described in Appendix 1of Kawamuro et al. (2016).
As the reference, we analyze the light-curve of the four
identiﬁed TDEs and ﬁnd that all of them satisfy the following two characteristics. The ﬁrst one is high variability amplitudes: the ratio between the highest ﬂux andthe one of the previous bin in the 30-day (90-day) averaged light curves is above 5. The second one is thatthe decay light curves are consistent with a power-law
proﬁle of t−5/3 , where t is time since the onset time ofeach TDE. We note that the time of the ﬂux peak, tp , isdelayed from the TDE onset time by approximately 5–80
days. Hence, we set the central day of the bin showing
the highest ﬂux as tp , and estimate the TDE onset timeby correcting for these oﬀsets. A power-law ﬁt to thelight curve over 90 days after the peak ﬂux is found tobe acceptable in terms of a χ2 test (Figure 1).
Accordingly, we apply the above two conditions to thelight curves of all MAXI sources (in total 506) except for
the four TDEs. As a result, we ﬁnd that none of themsatisfy the two criteria except for the objects identiﬁed as
active galactic nuclei (AGNs) or X-ray galactic sources.
Thus, we conclude that MAXI detected only the four
TDEs identiﬁed above during the ﬁrst 37 months of its
operation. This sample can be regarded as a statistically complete one at the sensitivity limit of MAXI fortransient events, as long as TDEs share similar characteristics in the X-ray light curve to those of the known

201

12

2
0

6
4
2

-2
55700

55750

55800

55850

MJD (DAY)NGC 4845
n=-2.00 (-1.25/+1.14)
6

6

8

0

-2
55650

Swift J2058+05
n=-1.92 (-0.73/+0.58)

8Flux (mCrab)

Flux (mCrab)

Flux (mCrab)

6
4

8
Swift J1644+57
n=-1.94 (-0.48/+0.39)

10

55600

55650

55700

55750

55800

Flux (mCrab)

Swift J1112-82
n=-1.86 (-1.55/+0.45)

8

4
2

4
2

0

0

-2

-2
55650

55700

MJD (DAY)

55750

55800

55850

55500

MJD (DAY)

55550

55600

55650

55700

MJD (DAY)

Fig. 1. X-ray (3–10 keV) light curves of four TDEs during their ﬂares. Only blue regions are used to be ﬁt with a power-law decay model.
The best-ﬁt power-law index (n) is indicated in each panel with errors at 90% conﬁdence level.

For TDEs with relativistic jets, the conversion factor
C1 can be written as

events.
3.

Hard X-ray Luminosity Function of Tidal Disruption
Events
3.1. Deﬁnition of Luminosities
In this section, we summarize the deﬁnitions of TDE
luminosities used in our analysis (Lobs
X and LX ). To makeluminosity conversion between diﬀerent energy bands, weneed to assume model spectra of TDEs. According to theprevious studies on TDEs observed in the X-ray band we
can approximate that the X-ray spectra of TDEs without
jets are composed of blackbody radiation and a powerlaw, which originate from the optically thick disk andits Comptonized component by hot corona, respectively.
For TDEs with jets, a relativistically beamed power law
is added to the above spectrum.
We deﬁne LX as the “intrinsic” peak luminosity of theComptonized power-law component in the rest-frame 4–
10 keV band. It can be converted into the “observed”
peak luminosity by
Lobs
X = CLX ,(1)

where C (C0 or C1 ; see below) is the conversion factor
that depends on the shape of the spectrum, redshift, andviewing angle with respect to the jet axis (for TDEs withjets).
Introducing the fraction of TDEs with jets in total
TDEs, fjet , we divide TDEs into two types, one withjets and the other without jets. For the TDEs without
jets, whose fraction is (1 − fjet ), the conversion factor C0
can be written as
C0 = ωpow + ωbb ,(2)

where ωpow and ωbb are those for the Comptonized(power-law) and blackbody components, respectively.
We normalize ωpow to unity when Lobs
of a TDE atX
z = 0 is deﬁned in the 4–10 keV band and its spectrum is
not absorbed. The factor ωbb depends on the broadband
spectrum of a TDE. Note that the blackbody component
is negligible in the 4–10 keV band because its temperature is expected to be much lower than a few keV.

C1 = ωpow + ωbb + ωjet ηjet δ 4(3)where the last term represents the jet contribution.
Here, ωjet takes account of the energy-band conversion,and ηjet is the fraction of the intrinsic luminosity (i.e.,that would be observed without beaming) of the jet in LX(the peak luminosity of the Comptonized component).
The Doppler factor (δ) is represented as

δ=

Γ−

√

1
Γ2 − 1 cos θ,(4)where Γ is the Lorentz factor and θ is the viewing angle
with respect to the jet axis. The observed luminosity
from the jet becomes larger than the intrinsic one by afactor of δ 4 with a frequency shift by δ. In our analysis,
we adopt Γ = 10, which is suggested from the analysisof the spectral energy distribution of Swift J1644+57
by Burrows et al. (2011). We also adopt ηjet = 0.1 asa standard value. The conversion factors C0 and C1
are dimensionless, composed only of the dimensionlessfactors (ω, fjet , and δ).
3.2. TDE Sample
We utilize the four identiﬁed TDEs listed in Table 1 as acomplete sample from our MAXI survey, and derive the
XLF. When integrated for 30 days, the MAXI survey
covers all the high Galactic latitudes (|b| > 10◦ ) region,
which corresponds to 83% of the entire sky. We havesearched for TDEs based on 30 days or 90 days binned
light curves. Thus, the sensitivity limit for the 30-days
averaged peak ﬂux of TDEs to which our survey is complete is determined by that for 30-days integrated data.
The resultant limit is 2.5 mCrab, or 3 × 10−11 erg s−1
cm−2 (4–10 keV).
3.3. Formulation of TDE X-ray Luminosity Function
We deﬁne the XLF of TDEs so that dΦ(LX , z)/dLX represents the TDE occurrence rate per unit co-moving volume per LX per unit rest-frame time, as a function of−1
LX and z, in units of Mpc−3 L−1
.
X yr

202

In our work, we make simple assumptions for modelling the shape of the TDE XLF. First, we write theTDE occurrence rate per unit volume as a function of
SMBH mass. It should be proportional to the product of the SMBH mass function (i.e., comoving number
density of SMBHs) and a speciﬁc TDE rate in a single SMBH. The local SMBH mass function can be derived from the local luminosity function of galaxies andthe Faber-Jackson relation between the galaxy luminosk
ity and the SMBH mass (Lgal ∝ MBH) as
ψ(MBH∗ ; MBH )dMBH = ψ0

 M

BH

MBH∗

γ −( MBH )k dM
BH
e MBH∗.
MBH∗(5)Here, γ = k(α + 1) − 1, where α is a parameter of thegalaxy luminosity function deﬁned as Ψ(Lgal∗ ; Lgal )dL =Ψ0 (Lgal /Lgal∗ )α e−Lgal /Lgal∗ dLgal /Lgal∗ . The subscript ∗indicates the characteristic parameter. Unless otherwise
noted, we adopt k = 0.8, Ψ0 = 0.007, log(MBH∗ /M ) =8.4, and α = −1.3 according to the results obtained by
Marconi & Hunt (2003) and Blanton et al. (2001), whereM is the solar mass. We refer to the dependence of thespeciﬁc TDE rate on SMBH mass derived by Stone &Metzger (2014),λξ ∝ MBH,(6)

where λ is chosen to be −0.4.
To represent the TDE occurrence rate as a functionof luminosity, we further make an assumption that thepeak luminosity L of a TDE (i.e., that free from the jetluminosity) is proportional to the SMBH mass, or equivalently, a constant fraction of the Eddington luminosity,
λEdd . Then, by converting MBH into L in the productof ψ(MBH∗ ; MBH ) and ξ, we can express the occurrence
rate of TDEs per unit volume in terms of L as
φ(L∗ ; L)dL = ψ0 ξ0 L γ+λL∗

for the mass of a SMBH that can cause a TDE. Hence,
the mass function of stars should be incorporated in calculating the actual TDE XLF. We approximate it by theshape of an initial mass function (IMF) with an upper
star mass boundary M∗,max , considering that very massive stars are already dead due to their short life time.
Here we employ the IMF derived by Kroupa (2001).
3.5. Maximum Likelihood Fit
We adopt the unbinned maximum likelihood (ML)method to constrain the XLF parameters. Whereas theML ﬁt gives the best-ﬁt parameters, the goodness of theﬁt cannot be evaluated. Hence, we perform one dimensional Kolmogorov-Smirnov test (hereafter KS test) separately for the redshift distribution and for the luminosity distribution between the observed data and best-ﬁt
model. The p-value, the chance of getting observed dataset, is evaluated from the one-sided KS test statistic.
We deﬁne the likelihood function as

L = −2i  ln     

where the subscript index i refers to each TDE and thedenominator represents the number of observable TDEs
with the intrinsic peak luminosity LX , the observed one
Lobs
X , the redshift z, the mass of the star M∗ , and theviewing angle θ, expected from the survey (note that θ
is related to the solid angle as dΩ = 2πd(cos(θ))). By
considering that the fraction of TDEs with jets among
all TDEs is fjet , the diﬀerential number is calculated as
obs

N (LX , LX , z, θ, M∗ )



obs
obs
= (1 − fjet )δD (C0 LX − LX ) + fjet δD C1 LX − LX×e−( LL∗ )k

dL
.
L∗

 ΔT
dΦ(LX , z) d2L (z) dτ  Lobs
XcA
P (M∗ ),
2dLX
1 + z dz
dL
1+z

(10)

(7)

We incorporate a redshift dependence of the TDE XLF
with an evolution factor of (1 + z)p that is multiplied to
the local XLF. Thus, the TDE XLF is formulated asdΦ(LX , z)
dLX = (1 + z)p φ(LX∗ ; LX )dLX .
dLX

N (LX , Lobs
Xi , zi , M∗ , θ)dLX d log M∗ dΩ/2π, (9)N (LX , Lobs , z, M∗ , θ)dLX dLobs dzd log M∗ dΩ/2π
X
X

(8)

3.4. Mass Function of Stars Disrupted by SMBHs
A TDE occurs only when the tidal disruption radius, RTDE = R∗ (MBH /M∗ )1/3 , where M∗ and R∗ are
the mass and radius of the star, is larger than the
Schwarzschild radius, RSch (≡ 2GMBH /c2 , where G and
c is the gravitational constant and the light speed, respectively.). For a given star, there is an upper boundary

where δD (x) is the Dirac δ-function, dL the luminosity
distance, dτ /dz the diﬀerential look-back time, A the2survey area at the ﬂux limit of Lobs
X /4πdL , and ΔT (=37 months) the survey time at the observer’s frame. The
factor 1/(1 + z) comes from the time dilation at z.
Since our sample size of TDEs is very small (four), we
ﬁx the following parameters of the XLF model, whichcannot be well constrained from the data. We adopt thecharacteristic luminosity of log LX∗ = 44.6 corresponding to the Eddington ratio of λEdd = 1 (see next paragraph). Table 2 lists the standard values of the Lorentz
factor of the jets, the dependence of the speciﬁc TDE
rate on SMBH mass, the fraction of the intrinsic jet luminosity in LX , and the upper mass boundary of tidally
disrupted stars. The index of the redshift evolution is assumed to be either p = 0 (no evolution case) or 4 (strong

203

Table 2. Default Setting of Fixed Parameters

Notes.
Col. [1]: Eddington ratio. Col. [2]: The Lorentz factor. Col.
[3]: Index for the MBH dependence of the TDE rate. Col.
[4]: Fraction of the intrinsic luminosity of the jet in LX . Col.
[5]: Upper mass boundary of disrupted stars in units of solar
mass.

evolution case). According to the prediction of numericalsimulations that the occurrence rate of TDEs increases
with the star-formation rate (Aharon et al. 2016), thelatter case simply assumes that the TDE rate is proportional to the star-formation rate density, which evolves
with ∝ (1+z)4 (Pérez-González et al. 2005). Eventually,
only φ0 ξ0 and fjet are left as free parameters.
We calculate the likelihood function over a redshift
range of z = 0–1.5. The luminosity range is derived
from the SMBH range of log(MBH /M ) = 4–8 for agiven λEdd (1.0). To convert the mass into the X-ray
luminosity LX , we consider λEdd = 1.0 and a spectra described in Section 3.1. (the blackbody, Comptonization,and jet components). Speciﬁcally, we ﬁrst determine theratio of the Comptonized component in the 2–10 keVband to the bolometric luminosity without the jet component. Assuming that accretion physics in TDEs issimilar to that of AGNs, we refer to the results by Vasudevan & Fabian (2007), who derived the bolometriccorrection factor (k2-10 ) from the 2–10 keV band to be∼50 for λEdd = 1.0. We then convert the luminosity
in the 2–10 keV band to that in the 4–10 keV band byassuming a power-law photon index of 2.0. As a result,log LX spans a range from 40.2 to 44.2 for λEdd = 1.0.
The integration range of M∗ is determined to satisfy thecriterion RTDE /RSch ≥ 1.
Table 3 summarizes the results of the ML ﬁt for thecosmological evolution index of p = 0 and p = 4. The
results are acceptable in terms of the KS-test at 90%
conﬁdence level. The fjet upper and lower limits, where
the L-value is increased by 2.7 from its minimum, arederived. Since the ML method cannot directly determine the normalization (ψ0 ξ0 ) of the luminosity function, we calculate it so that the predicted number fromthe model equals to the detected number of the TDEs.
The attached error corresponds to the Poisson error in
the observed number at the 90% conﬁdence level basedon equations (9) and (12) in Gehrels (1986).
Figure 2 displays the results of the TDE XLF as a function of Lobs
X (observed peak luminosity in the 4–10 keVband), obtained for the case of p = 4. The solid curves

p-value (Lobs
X -dist/z-dist)[4]
0.16/0.14
0.60/0.40

ψ0 ξ 0[3]
1.7+1.6−0.9
1.6+1.6−0.9

Notes.
Col. [1]: Index of the redshift evolution. Col. [2]: Fraction of
the TDEs with jets. Col. [3]: Normalization factor of XLF in−1 . Col. [4]: p-value on theunits of 10−8 Mpc−3 log L−1
X yrbasis of the KS-test for each parameter of Lobs
X and z.

plot the best-ﬁt model at z = 0.75, which is obtainedby integrating dΦ(LX , z = 0.75)/dLX over the half solid
angle with respect to the jet direction, the mass of disrupted stars, and LX (intrinsic peak luminosity of theComptonized component) as
obs

Π(log LX , z = 0.75)
  
dΩ
obs
dLX d log M∗= ln(10) LX
2π



obs
obs
× (1 − fjet )δD (C0 LX − LX ) + fjet δD C1 LX − LX
×dΦ(LX , z = 0.75)
P (M∗ ).
dLX

(11)

The data points are plotted by the “N obj /N model ”
method (Miyaji et al. 2001). The error bars reﬂecting
the 90% conﬁdence level of N obj are derived based on
Gehrels (1986).
p=40.0001-1M∗,max
[5]
1.0

1e-06

-3

ηjet
[4]
0.1

Mpc yr ]

λ
[3]
-0.4

obs -1

Γ
[2]
10

fjet
[2]
0.012+0.122−0.0100.003+0.027−0.002

Π [log LX

λEdd
[1]
1.0

Table 3. Best-ﬁt parameters

p
[1]
0
4

1e-08

1e-10

1e-12

1e-14
42

43

44

45
obs

log LX

46

47

48

49

-1

(4-10 keV) [erg s ]Fig. 2. The best-ﬁt XLF as a function of “observed” peak luminosity
at z = 0.75 for an evolution index of p = 4. The solid line
represents the total XLF consisting of that of TDEs without jets
(dotted line) and that of TDEs with jets (dot-dashed line).

3.6. Mass Accretion History of SMBHs by TDEs
TDEs contribute to the growth of SMBHs as argued by
Soltan (1982) for AGNs. We here calculate the evolution of the mass density of SMBHs by TDEs as done
for AGNs. The bolometric luminosity of a TDE can be
related to the mass accretion rate Ṁacc via the mass-toradiation conversion eﬃciency  as
Lbol = c2 Ṁacc .

(12)

204

The mass growth rate of the SMBH is given by
ṀBH = (1 − )Ṁacc .

(13)

We formulate the SMBH mass-density equation as
ρ(z) =

z

zs

dt
dz dz

× 1−
c2

∞
tp

 LX,max
LX,mindLXdΦ(LX ,z)  M∗,max
M∗,min P (M∗ )d log M∗
dLX{(1 − fjet )Lpeak + fjet Ljet
peak }

t
tp

−5/3

dt,

(14)

where zs , Lpeak , and Ljet
peak are the initial redshift fromwhich the calculation starts, the peak bolometric luminosity without jets, and that with jets. Here, Lpeak islog ρ(z)/105 [Mo• Mpc-3]

set to the Eddington luminosity, while Ljet
peak is the sumof the Eddington luminosity and intrinsic jet luminosity.
By adopting the mass-to-radiation conversion eﬃciencyof  = 0.1 similarly to the case of AGNs, we calculate
the cumulative SMBH mass-density (M Mpc−3 ) fromzs = 1.5 for the case of p = 4. We ﬁnd that the total massdensity at z = 0 would become at most 7×102 M Mpc−3
by TDEs only, with a very small increase (3.8 × 105 M
Mpc−3 ) from that at z = 1.5 (Figure 3). Figure 3 compares the cumulative SMBH mass-density by TDEs andthat of AGNs estimated by Ueda et al. (2014), indicatingthat the mass-density contributed by TDEs is much lessthan that of AGNs.

1
0.9
0.8
0.7
0.6
0 0.2 0.4 0.6 0.8 1 1.2 1.4
zFig. 3. The evolution of SMBH mass-density caused by TDEs (red)and AGNs (black).

4. Summary
We have derived the XLF of TDEs, i.e., the occurrence
rate of a TDE per unit volume as a function of intrinsic
peak luminosity, from the MAXI extragalactic survey.
Our sample consists of four TDEs detected in the ﬁrst
37 months MAXI /GSC data at high Galactic latitudes(|b| > 10◦ ). It is complete to a ﬂux limit of ∼ 3 × 10−11
erg s−1 cm−2 (4–10 keV).
We formulate the shape of the TDE XLF, based on
the mass function of SMBHs and the speciﬁc TDE rate

as a function of SMBH mass. We take account of two
distinct types of TDEs, those with jets and those withoutthem, and also the relativistic beaming from the jets.
To incorporate eﬀects of the cosmological evolution, we
assume two cases where the XLF is constant over redshift
or is proportional to (1 + z)4 . ML ﬁts are performed tothe observed TDE sample, with the normalization of the
XLF (i.e., TDE rate) and the fraction of TDEs with jets
among all TDEs, fjet , allowed to vary. Resultant fjet isconstrained to be 0.0007%–34%, consistent with the caseof AGNs. Eventually, based on the XLFs, we ﬁnd thatthe eﬀect by TDEs to the growth of SMBHs is negligible
at z  1.5.
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1235

205

First Statistical Tests for Clumpy-Torus Models:Constraints from RXTE Monitoring of Seyfert AGN
Alex Markowitz,1 Mirko Krumpe,2 and Robert Nikutta3
1

University of California, San Diego, Center for Astrophysics and Space Sciences, MC 0424,
La Jolla, CA, 92093-0424, USA
2
Leibniz-Institut für Astrophysik Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany3
National Optical Astronomy Observatory, 950 N Cherry Ave., Tucson, AZ, 85719, USA
E-mail(AM): almarkowitz@ucsd.edu
Abstract
This presentation is a summary of highlights from Markowitz, Krumpe, & Nikutta (2014) and preliminary results from our follow-up paper, Nikutta, Krumpe, & Markowitz, in prep. We present an analysis ofmulti-timescale variability in line-of-sight X-ray absorbing gas as a function of optical classiﬁcation in AGN
to derive the ﬁrst statistical constraints for recent clumpy absorbing torus models. Such models represent
the paradigm shift away from the classical assumed solid donut morphology and towards a morphology
containing numerous discrete clouds or clumps. Such advances come courtesy of sustained long-term X-ray
monitoring; we use the vast archive of Rossi X-ray Timing Explorer (RXTE ) multi-timescale monitoringof dozens of type I and Compton-thin type II Seyfert AGN. We search for discrete absorption events due
to clouds transiting the line of sight; most of our twelve detected clouds are Compton-thin and are locatedin the outer BLR or inner dusty torus. We discuss the resulting implications for cloud distributions in thecontext of Clumpy torus models. We discuss cloud sizes, stability, and radial distribution across a wide
range of distances, and explore the exhibited range in density proﬁles for the highest-quality eclipse events.
We discuss possible connections to the mechanisms that form and launch clouds. In addition, all observed
clouds are sub-critical with respect to tidal disruption; self-gravity alone cannot contain them. External
forces, such as magnetic ﬁelds or ambient pressure, are needed to contain them. Otherwise, clouds must
be short-lived. Our results apply to both dusty and non-dusty clumpy media, and probe model parameter
space complementary to that for short-term eclipses observed with XMM-Newton, Suzaku, and Chandra.
Key words: galaxies: active – X-rays: galaxies – galaxies: Seyfert

1.

Introduction

The exact morphology of the circumnuclear gas in active
galactic nuclei (AGN) remains unclear, but holds the key
to answering how material gets funneled from radii of kpc
in the host galaxy down to the accretion disk, and ultimately feeds the supermassive black hole. This unsolved
question impacts our understanding of the eﬃciency of
black hole feeding as well as AGN duty cycles.
The community has long been in the process of shifting
away from quantifying emission and absorption processes
in Seyferts by modeling circumnuclear gas via a simpliﬁed homogeneous, Compton-thick donut morphology.
Instead, a new generation of models describe the torus
via distributions of numerous individual clouds (Elitzur
2007; Nenkova et al. 2008), although suggestions thatthe torus should consist of clouds go as far back as, e.g.,
Krolik & Begelman (1986); Krolik & Begelman (1988).

Often, the clouds are embedded in some outﬂowing windfrom the cold, thin accretion disk which feeds the black
hole. Observational support for these models so far hascome mainly from ﬁtting IR SEDs in small samples of
Seyferts (Ramos Almeida et al. 2011; Alonso-Herrero etal. 2011).
Evidence for clumpy tori also comes from variations
in column densities NH of line-of-sight X-ray absorbing
columns in both (optically classiﬁed) type Is and IIs,
with timescales of variability ranging from hours to years(Risaliti, Elvis, & Nicastro 2002). The community has
also observed moderately Compton-thick variations in
NGC 1365 (Risaliti et al. 2007; Risaliti et al. 2009) andNGC 7582 (Bianchi et al. 2009), as well as rapid changes
in covering fractions of partial-covering absorbers (Puccetti et al. 2007; Risaliti et al. 2011; Sanfrutos et al.
2013). When adequate data exist, time-resolved spec-

206

troscopy of full eclipse events has yielded constraints
on density proﬁles in the transverse direction for longduration (3–6 months) eclipses in NGC 3227 (Lamer et
al. 2003) and Cen A (Rivers et al. 2011) and for eclipses∼1 d in NGC 1365 (Maiolino et al. 2010) and SWIFT
J2127.4+5654 (Sanfrutos et al. 2013). In addition, X-ray
obscuration is the only way to probe obscuration insidethe dust sublimation radius Rd , and to test the notionthat the dust-free BLR and dusty torus form a common
radially extended structure inside and outside Rd , resepctively (Netzer & Laor 1993; Elitzur 2007; Gaskell et
al. 2008).
Typically for clumpy-torus models such as the
Clumpy models by Nenkova et al. (2008), the cloud distribution is preferentailly concentrated towards the equatorial plane, and is described by a Gaussian angular distribution σ, mean number of clouds along an equatorial
ray N0 , and a given radial power-law distribution out tosome maximum radius, with each cloud being identicaland having the same optical depth.
A statistical survey of the environments around supermassive black holes has been needed to properly constrain parameter space in the Clumpy models. These
parameters so far have been constrained only from IR
SED ﬁtting, as mentioned above. In our 2014 paper, Markowitz, Krumpe, & Nikutta (2014; hereafter
MKN14), we present the ﬁrst such X-ray-absorptionbased survey, the longest AGN X-ray monitoring study
to date. Our survey quantiﬁes line-of-sight X-ray absorption by clouds that transit the line of sight to thecentral engine. We assessed the relevance of Clumpy
torus models as a function of optical classiﬁcation by
exploring absorption over a wide range of length scales(both inside and outside the dust sublimation zone).
2. Methodology
We consider archival data collected with RXTE ’s Proportional Counter Array (PCA). RXTE was the ﬁrst Xray mission to operate using sustained monitoring campaigns, with regularly spaced visits, usually 12 ks each,over durations of weeks to years.
RXTE visited 153 AGN over its lifetime. We only
consider sources with mean 2–10 keV ﬂux  8 × 10−12
erg cm−2 s−1 , and visited at least four times during themission. We do not consider blazars. Due to the PCA’s
energy resolution, we cannot accurately quantify changes
in NH when a steady full-covering Compton-thick absorber is present, other than verifying that none of theseven Compton-thick sources monitored by RXTE transitioned to/from being Compton-thin. The ﬁnal target
list consists of 37 type I and 18 Compton-thin type II
AGN. Here, we classify “type I” to include Sy 1.0, 1.2,
and 1.5; we classify “type II” to include Sy 1.9 and 2(we had no type 1.8s in our sample). 11/18 type IIs

in our sample show evidence for “hidden” BLRs, manifested via scattered/polarized optical emission or via IR
Paschen lines, and we adhere to the assumption that
BLRs do exist in all our objects, and that optical classiﬁcations are due to increasing levels of obscuration in the
optical band, independent of assumptions about systemorientation.
We examined 7–10/2–4 keV hardness ratio (HR1)light curves; assuming full-covering absorption, HR1
peaks at column densities roughly 1 − 3 × 1023 cm−2
while giving us sensitivity down to ∼ a few 1022 cm−2 .
A simple power-law ﬁt to the 2–10 keV band, not accounting for absorption, can also indicate excess absorption: measurements of the “apparent” photon index suddenly falling to values lower than ∼ 1.5 − 1.6, the typically ﬂattest coronal power-laws measured in unabsorbed
Seyferts, are likely due to absorption events. We distinguish between “secure” vs “candidate” eclipse events:
the former group are those conﬁrmed via binned timeresolved spectral ﬁtting to have a temporary increase in
NH . Additional analysis details and eclipse identiﬁcation
criteria are presented in MKN14.
As a caveat, we are only able to detect an absorption
event and rule out spectral pivoting if the cloud’s covering fraction is greater than ∼ 80 − 90%, and if its ionization parameter ξ [erg cm s−1 ] is less than log ξ ∼ 1 − 2.
Our ﬁndings are thus complementary to those derived
using the archival databases of other X-ray missions:
XMM-Newton, Chandra, and Suzaku can probe clouds
with lower column densities, partial covering, and/or
high-ionization absorption, and eclipse events with durations 1–3 d (Turner et al. 2008; Risaliti et al. 2011).
However, unlike RXTE, these missions generally do not
perform sustained monitoring campaigns and are notable to detect absorption events longer than ∼ a few
days. In addition, RXTE -PCA’s coverage above 10 keV
aﬀords increased sensitivity to higher absorbing columns.
3. Summary of main results
We ﬁnd twelve “secure” X-ray absorption events (conﬁrmed with spectral ﬁtting) across eight Seyferts:
NGC 3783, NGC 3227, MR 2251–178, Mkn 79, Mkn 509,
Cen A, Mkn 348, and NGC 5506. Four of theseeclipse events were reported previously (Lamer et al.
2003; Rivers et al. 2011; Rothschild et al. 2011; Smithet al. 2001; Akylas et al. 2002). We also report an
additional four “candidate” events in three Seyferts(NGC 3783, NGC 3516, and Fairall 9).
The “secure” events span durations from less than one
day (for NGC 5506) to 1.5 years (for Mkn 348). The
clouds are modeled to have column densities spanning
4 − 26 × 1022 cm−2 ; we did not detect any full-covering
Compton-thick clouds. Assuming a Galactic gas/dust
ratio, these X-ray columns imply dust with V-band ex-

4. Resolving clouds’ density proﬁlesWhen adequate time-resolved data are present, we canobtain information on a given cloud’s density proﬁle in
the transverse direction, and recent results are beginning
to reveal a diversity in cloud proﬁle shapes.

5

10

NGC 3783

0

tinction AV = 22 − 144 mag., and V-band optical depths
τV of 20–132. These values are consistent with optical depths typically used in Clumpy modeling (Nenkova
et al. 2008) and derived by ﬁts to IR SEDs by AlonsoHerrero et al. (2011) and Ramos Almeida et al. (2011).
To estimate the distance to a given cloud from the supermassive black hole rcl , we use the measured eclipse
duration and column density, combined with constraints
on the cloud ionization level. Speciﬁcally, we follow §4 of
Lamer et al. (2003) and assume that clouds are in Keplerian orbits and that each cloud has a uniform density.
Cloud diameter Dcl = NH /nH = vcl tD , where tD is themeasuredeclipse time and vcl is the transverse velocity,
equal to GMBH /rcl . Using the deﬁnition of ionization
parameter ξ, one obtains (see Eqn. 3 of Lamer et al. 200
1/5 2/5 2/5 −2/53): rcl = 4 × 1016 M7 L42 tD NH,22 ξ −2/5 cm, where
M7 = MBH /106 M , L42 = Lion /(1042 erg s−1 ), NH,22 =NH /(1022 cm−2 ), and tD is in units of days.
Our constraints on the ionization parameters for ourclouds are poor. Given the column densities of our clouds
and the energy resolution of the PCA, we can safely ruleout values of log(ξ) above ∼ 2. We therefore calculatedistances assuming log(ξ) = −1, 0, or +1. Uncertainty
on rcl is thus dominated by uncertainty on ξ.
We obtain best-estimate values of rcl that are typicallytens to hundreds of light-days from the central engine,and lie in the range 1 − 50 × 104 Rg (0.3 − 140 × 104 Rg
accounting for uncertainties).
The clouds in 7/8 objects (all except NGC 5506) areconsistent with the dust sublimation zone considering
uncertainties on rcl , i.e., in the outer BLR and/or innerdusty torus. However, this “clustering” may be in partassociated with our observational bias to select eclipses
with events of ∼tens of days, as per our selection function. In Cen A, the clouds are inferred to be consistent
with residing entirely in the dusty zone. In contrast,
NGC 5506 has the lowest value of rcl /Rd ; this cloud islikely the least dusty of the secure events in our sample,
and is likely commensurate with that source’s BLR.
The average cloud diameter was ∼ 5 light-hours(roughly the Sun–Neptune distance). In addition, theinferred diameters and locations for these (full-covering)
clouds imply upper limits to the size of the X-ray
continuum-emitting region. These limits are as low as
2Rg for NGC 5506 and 45−65Rg for NGC 3783, Mkn 79,
and Mkn 509. The average cloud mass was 10−5.2 M ,and the average number density was 108.5 cm−3 .

NH (1022 cm−2)

207

5

10

3

25

30

35

2

]


20

Days since MJD 54550

wift{ Suzaku

3


15

mc 


[ 1
H
N

5
0

54770

54780
54790
54800
d e (Mod ed J l n Dy)Fig. 1. The community is starting to amass a variety in behavior inclouds’ column density proﬁles as traced by NH (t), when adequate
sampling and signal-to-noise allow. (Top panel) — From MKN14:
an 18-day long eclipse event in NGC 3783 as traced by RXTE,
revealing two column density peaks. (Bottom panel) — From
Beuchert et al. (2015): a ﬁve-week Suzaku/Swift campaign on
NGC 3227 in 2008 caught the source obsured by an ionized cloudwith an irregular density proﬁle.

For example, the 2010–11 eclipse in Cen A observed byRivers et al. (2011) had a NH proﬁle that was symmetricin time and centrally-peaked; Rivers et al. (2011) foundthat a centrally-peaked number density ﬁt better thana uniform number-density sphere. Lamer et al. (2003)
and Sanfrutos et al. (2013) also observed symmetric NH
proﬁles in the 2000–1 eclipse of NGC 3227 and an eclipse
in SWIFT J2127.4+5654, respectively.
However, many other eclipse proﬁles reveal embedded overdense/underdense regions: a double-peaked NH
proﬁle in NGC 3783 observed by MKN14, a highlyvariable proﬁle in a NGC 3227 event during a ﬁveweek Suzaku/Swift campaign measured by Beuchert etal. (2015) (see Fig. 1), and “comet” and “anti-comet”
shapes observed in NGC 1365 by Maiolino et al. (2010)
and Rivers et al. (2015).
These proﬁles may be clues as to how clouds areformed and sculpted, e.g., the torus may be a dynamically active place, with eclipsing clouds possibly being due to MHD winds launched upwards from thedisk (Fukumura et al. 2010; Czerny & Hryniewicz 2011)
and/or compressed by poloidal magnetic ﬁelds (Emmering et al. 1992).

208

5. On non-variable absorption in type IIs
We ﬁnd evidence in eight type II Seyferts for a baseline
level of X-ray absorption that is constant in time overtimescales from 0.6 to 8.4 years; the observed hardness
ratios place limits on ΔNH ∼ 0.6 − 9 × 1022 cm−2 (seeFig. 2). In the context of clumpy-torus models explaining all the observed absorption, one would require that alarge number of very low-column density ( 1022 cm−2 )clouds along the line of sight, with the total number ofclouds along thre line of sight remaining roughly constant within each object. Alternately, the constant levelof absorption can be explained by absorbing structuresassociated with the host galaxy, e.g., dust lanes (Gould
et al. 2012). Another possibility is a medium of nonclumpy, highly homogeneous gas close to the black hole,such as a homogeneous intercloud medium.
However, in Cen A, MKN14 detected an 80-day “antieclipse”: the baseline level of absorption temporarily
dipped by ∼ 14% then recovered. This event suggeststhat material close to the black hole is not highly homogeneous, and that we witnessed a transit across the lineof sight by a relatively underdense region.
6. Instantaneous probabilities to witness an eclipse
We derived the probability to catch a type I/II source undergoing an eclipse event that has any duration between
0.2 d and 16 yr, taking into account the highly inhomogeneous sampling in our X-ray observations. That is, weestimated the probability to instantaneously observe asource undergoing an eclipse event due only to a cloud
passage through the line of sight, and independent oflong-term constant absorption, e.g., associated with gas
in the host galaxy.
For type Is, it is 0.006 (error range: 0.003–0.166); fortype IIs, 0.110 (0.039–0.571). Our uncertainties are conservative, as they take into account our selection function, candidate eclipse events in addition to the secureones, uncertainties in the observed durations, and uncertainties in the contributions of individual objects sampling patterns to our total sensitivity function; we useda Monte Carlo ”bootstrap” method to estimate uncertainties.
Although subject to low number statistics, our observations hint at diﬀerences in the distributions of observed
eclipse event durations and probabilities for type I andII objects. In addition, assuming that all type Is or
IIs have a common cloud distribution morphology, thenthese instantaneous probability values can be comparedto Clumpy theoretical predictions that assume certain
cloud distribution parameters (Clumpy parameters inclination angle, angular distibution of clouds, and number of clouds along an equatorial ray) to provide constraints on these parameters.
However, as a caveat, these probabilities were derived

with a sample that had some biases due to highly inhomogeneous X-ray sampling, and due to type Is’ being
sampled ∼ 4 − 5 times more often than type IIs. We
were also limited in sensitivity to changes in column density no lower than ∼ 1022 cm−2 . In addition, we were
not highly sensitive to absorption by highly-ionized or
partial covering gas; our PCA spectra were sensitive toionization levels up to log ξ ∼ 1 − 2. When one considers
the full range of possible cloud properties, the resulting
probabilities will almost certainly be higher.
7. Preliminary results from Nikutta et al., in prep: additional constraints on cloud properties
We also present a brief overview of results from our forthcoming paper in which we derive additional constraints
on cloud properties.
The “long-term event” clouds in MKN14 inferred toreside in the outer BLR/inner dusty torus and the“short-term event” clouds reported by Risaliti et al.
(2007), Risaliti et al. (2009), Maiolino et al. (2010), Sanfrutos et al. (2013), and Turner et al. (2008) and inferred
to reside in the BLR are all roughly consistent with acommon angular size as seen from the central source,∼ 1.4 arcmin. That is, cloud physical diameters are consistent with increasing linearly with linear distance fromthe central black hole. This fact may provide a hint atcloud formation mechanisms: if clouds are formed fromthe disk, for example, various cloud fragmentation mechanisms e.g., Amaro-Seoane & Chen (2014), may be applicable.
This small average angular size also suggests thatin order for clumpy-torus models to explain the total
observed median line of sight absorption in type I/II
Seyferts (Ramos Almeida et al. 2011), a total of 4 × 107(type Is) or 8 × 107 (type IIs) clouds may exist across the
BLR + dusty torus (assuming no cloud overlap). These
values assume that clouds only account for all absorption; the presence of an intercloud medium or absorption due to host galaxy dust lanes along the line of sightwould of course yield lower cloud numbers.
We also infer that the 12 secure clouds events from
MKN14 do not possess suﬃcient self-gravitational resistance to tidal shearing. A cloud in orbit will be able
to withstand gravitational torque only if its density isgreater than the “ambient” local density of the black holemass spread out evenly over the sphere of the cloud’s
orbit. We use the estimates of cloud number densityfrom MKN14, and ﬁnd that clouds are underdense byfactors of 10–1000 to resist tidal shearing. The ambient pressure from an inter-cloud medium could suﬃce toconﬁne clouds (Netzer 2013). Alternatively, an ambient
magnetic ﬁeld can conﬁne clouds, as suggested by Rees
(1987) for BLR clouds (although the argument also applies to torus clouds); we estimate that the minimum re-209quired magnetic ﬁeld strength is only a few milli-Gauss.
8. Future prospects
While the RXTE sample represents the best current opportunity to explore long-term eclipses and clouds inthe outer BLR/dusty torus, we were limited to exploring neutral or near-neutral, full-covering or near-fullcovering (covering fraction  90%) clouds.
To eﬀectively explore larger regions of [ NH , covering fraction, ξ, Δt ] parameter space in the future,
our idealized ”X-ray monitoring wish list” for obtaining additional constraints from X-ray eclipses thus includes: •Sustained monitoring with evenly-spaced sampling, down to timescales of hours, with consistentsource-to-source coverage to mitigate observing biases
as a function of time. •Soft X-ray coverage to accesslower column densities and lower values of ΔNH compared to PCA, and to access constraints on cloud ionization parameter. •Access to lower 2–10 keV ﬂuxes
and a larger number of type IIs, as they are expected
to harbor a higher number of line of sight clouds iforientation is the dominant factor in type I/II uniﬁcation. In the context of clouds being produced in adisk wind, however, BLR/torus cloud production may
be hampered towards low luminosities (Elitzur & Ho2009), and cloud numbers as a function of luminosity
across the BLR and dusty torus must be tested. We alsoneed to test whether “baseline” levels of absorption are
truly constant, indicating either host galaxy absorption
or a highly-homogeneous intercloud medium, or variable,suggesting transiting over- or under-dense regions close
to the AGN. •We must span a wider range of luminosities than in MKN14 to test if covering factor evolves withluminosity, as implied by multiple high-redshift studies(Ueda et al. 2014). •At least CCD-quality spectral monitoring to fully deconvolve variations in NH and Γ, andconﬁrm partial-covering events. •Ideally, if optical/UV
monitoring is in place simultaneous to X-ray monitoring,such monitoring can tell us whether the cloud is dusty.
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210

NGC 5506
1998

1999

2000

2001

2002

2003

2004

5 10 15 20

1997

2
1
1.5
1
0.5

F7−10//F2−4

F2−4/μ F7−10/μ

3

F2−10

1996

5.05×104

5.1×104

5.15×104

5.2×104

5.25×104

5.3×104

Fig. 2. From MKN14: RXTE -PCA light curve of the Sy 1.9 NGC 5506, showing light curves for the 2–10, 2–4, 7–10 keV bands, and the
7–10/2–4 keV hardness ratio. The hardness ratio is consistent with a baseline level of absorption that is constant down to ΔNH ∼ 1× 1022
cm−2 for a duration of eight years; the exception is the short-term cloud transit in 2000. (The mild shift in average harness ratio in 1999
is instrumental, and not intrinsic to the source.)

211

A Novel View of AGN Accretion Flows Revealedby X-ray and Optical Monitoring
Kazuo Makishima,1 Hirofumi Noda,2 Takeo Minezaki,3 Mitsuru Kokubo,3 Tomoki Morokuma,3 Mamoru Doi,3
Yuki Moritani,3 Kazuhiro Nakazawa,3 Makoto Watanabe,4 Hikaru Nakao,5 Masataka Imai,5 Kenji Kawaguchi,6Katsutoshi Takagi,6 Koji S. Kawabata,6 Tatsuya Nakaoka,6 Miho Kawabata,6 Makoto Uemura,6Yasushi Fukazawa,6 Michitoshi Yoshida,6 Kumiko Morihana,7 Yoich Itoh,7
Yuhei Takagi,8 Akira Arai,9 Ryosuke Itoh,10 Yoshihiko Saito,10 and Shin’ya Yamada11
1

RIKEN; 2 Tohoku University; 3 University of Tokyo; 4 Okayama University of Science; 5 Hokkaido University;
6
Hiroshima University; 7 University of Hyogo; 8 National Astronomical Observatory of Japan;
9
Kyoto Sangyo University; 10 Tokyo Institute of Technology; 11 Tokyo Metropolitan University
E-mail(FS): maxima@riken.jp
Abstract
The Seyfert galaxy NGC 3516 was monitored for about a year, with Suzaku and ﬁve Japanese opticaltelescopes (Pirka, Kiso Schmidt, Nayuta, MITSuME, and Kanata). In both X-rays and optical, NGC 3516
remained in its very faint state. The source meantime varied in X-rays by a factor of ∼ 20, and in the
optical B-band by a similar amount in the ﬂux diﬀerence. A tight correlation was detected between thetwo bands, with an optical delay by 2.0+0.6−0.7 days. The optical change is hence due to X-ray reprocessing,but the lag is too large for the standard “lamppost” X-ray source illuminating an accretion disk that is
not truncated. The results are successfully explained by a hot accretion ﬂow inside a truncated disk.
Key words: AGN: X-rays — AGN: optical — AGN: individual: NGC 3516 — black holes — accretion

1. Introduction
One of the central issues in extragalactic astrophysics has
been the spectral energy distribution from active galactic
nuclei (AGNs), since it provides a vital clue to thesemass-accreting massive black holes (BHs) (Rees 1984).
There, major roles have been played by the X-ray andoptical bands, which bracket the ultra-violet range where
the emission is strongest but absorbed.
So far, broadband X-ray and optical signals from a
disk-dominated AGN has usually been interpreted in thefollowing scenario (e.g., Miniutti et al. 2007):
1. The matter accretes through an optically-thick standard accretion disk, which continues down to thelast circular orbit (ISCO) around the central BH.
2. A “lamp-post” type source, located right above the
BH, Comptonizes the disk photons into the powerlaw (PL) shaped hard X-ray continuum.
3. The hard X-rays illuminate innermost disk regions
where general relativity sets in, and produce asmeared reﬂection continuum and a broad Fe-K line.
4. The optical signals arise from the disk, either spontaneously, or by reprocessing the X-rays.
Even putting aside the reality of the lamp-post hard

X-ray sources, the above scenario, though popular, encounters two (or more) problems. One is the implicit
assumption, that the primary X-ray continuum consistsof a single PL component, is not fully tested againstcomplex AGN spectra which often allow diﬀerent interpretations (Cerruti et al. 2011). The other is that the
expected tight optical vs. X-ray correlation was observed
only occasionally (e.g. Maoz et al. 2002). Thus, theabove lamp-post scenario needs a critical revision.
Through Suzaku observations of MCG 6–30-15, NGC
3516, NGC 3227 (Noda et al.2011,2013,2014), IC 4329A(Miyake et al.2016,2017) and NGC 4051 (Seino al. 2017),
their 2–45 keV emission has been found to generally consist of two distinct primary continua. One is Soft and
Rapidly-varying Primary Component (SRPC) with aphoton index of Γ ∼ 2.2, and the other is Hard andGradually-varying Promary Component (HGPC) withΓ ∼ 1.5. The HGPC is observed throughout, whereas
the SRPC emerges only when the luminosity exceeds∼ 0.1% of the Eddington limit, LEdd . This discovery
answers negatively to the ﬁrst problem raised above, andpossibly solve the second problem too, because the optical signals can be correlated only with either the SRPCor HGPC, which were not distinguished previously.

212

3. Data Analysis and Results
3.1. X-ray light curves
Figure 1 shows two-band XIS light curves of NGC 3516
in the 7 observations. The source varied by a factor of
20 over the year, with some intra-day variations.
The source was rather X-ray faint throughout, withthe unabsorbed 2–10 keV luminosity (at 41 Mpc) of
L2-10 = (0.2 − 5.3) × 1042 erg s−1 . Employing the AGN
mass of MBH = 3.2 × 107 M (Denney et al. 2010) anda bolometric correction factor of 10, these translate to
Eddington ratios of η ≡ L2-10 /LEdd = (0.4 − 11) × 10−3 .

Count rate (cnt/s)

(a)

(b)

(c)

(d)

(e)

(f)

(g)

1
3−10 keV
0.1

2−3 keV
0.01
0

40

80

0

40 0

40

80

0

40

80

0

40

80

0

40

80 0

40

80

Time (ksec)

Fig. 1. Suzaku XIS light curves of NGC 3516 obtained on the 7 observations, in 2–3 keV (open circles) and 3–10 keV (ﬁlled circles).
Each data point is 5 ksec. From Noda et al. (2016).

1

−1

keV )

0.01

s

−1

0.01

−2

-3

10

keV (Photons cm

Photons cm

−2

s

−1

keV

−1

0.1

0.1

-4

10

10

-3

2

3
χ

2. Observations
To better understand the X-ray vs. optical correlation
in AGNs under the new light of the two distinct primarycontinua, we conducted extensive simultaneous observations of NGC 3516. This Type 1.5 Seyfert was chosenbecause it clearly showed the co-existence of the SRPCand HGPC (Noda et al. 2013), and its high declination,
δ = 72d 34m , makes it easily accessible with Northernhemisphere ground telescopes.
In the Suzaku AO-8 cycle, seven X-ray observationsof NGC 3516 were conducted from 2013 April to 2014
April, with various intervals from 1 week to 5 months.
Each of them had an exposure of ∼ 50 ksec, so that
we can utilized the new variability-assisted spectroscopy
technique called C3PO method (Noda et al. 2013; Noda
et al. 2014). We invested the ﬁrst ﬁve observations onto
a 50-days period from 2013 April 9 to May 29. Further
details are available in Noda et al. (2016).
The optical photometric monitoring of NGC 3516 was
carried out with 5 ground-based telescopes in Japan.
They are the Pirka (Nayoro), MITSuME (Akeno),
Kiso Schmidt, Nayuta (Nishi-harima), and Kanata (Hiroshima) telescopes, with a diameter of 1.6, 0.5, 1.5, 2.0,
and 1.5 meters, respectively. The photometry was carried out mainly in the B and V bands, using CCD detectors located at respective focal planes. With thesetelescopes, we observed NGC 3516 not only during theseven Suzaku observations, but also on other periods, as
frequently as possible, even without simultaneous X-ray
coverage. Details are also found in Noda et al. (2016).

0
-3
2

5

10
Energy (keV)

40

10

-4

2

5

10
Energy (keV)

40

Fig. 2. (left) The 2–45 keV Suzaku XIS and HXD spectra of NGC
3516 from the 7 observations, shown after background subtraction
but without removing the instrumental responses. They are ﬁtted
simultaneously with a cutoﬀ-PL plus reﬂection model. (right) The
inferred best-ﬁt models for the 7 data sets, shown in νF ν form.
Taken from Noda et al. (2016).

3.2. X-ray spectra
The XIS and HXD spectra of NGC 3516 from the 7 observations are presented in Fig. 2 (left). We ﬁtted themjointly, with a model consisting of a cutoﬀ-PL continuum
and its cold reﬂection (including the Fe-K line), modiﬁed by photo-absorption. While the PL normalization
was allowed to diﬀer, Γ was constrained to be common
among the 7 epochs. The reﬂection component was assumed to have a 60◦ inclination, an Fe abundance of 1.0
Solar, and a normalization that depends on the epochs.
As shown in Fig. 2 (left), the model successfully reproduce the 7 spectra, with χ2 /dof = 1184.0/1113. The PL
slope was obtained as Γ = 1.75±0.02, and the data show
no evidence of broad Fe-K line. The best-ﬁt models are
drawn in Fig. 2 (right).
In spite of the large continuum variation, the reﬂection
varied only by a factor of 2 (see Fig. 2). As a result,
the reﬂector’s nominal solid angle was consistent with
2π in the 2nd – 5th data sets, whereas it increased to(4 − 6) × 2π when the source was dimmer. The lattercan be understood as afterglows of the distant reﬂector.
In the 2005 Suzaku observation when the unabsorbed
2–10 keV luminosity was relatively high at 5.9 × 1042 erg
s−1 (Noda et al. 2013), NGC 3516 exhibited a variable
PL component with Γ ∼ 2.2 identiﬁed with the SRPC, astable PL with Γ ∼ 1.1 interpreted as the HGPC, and acold reﬂection with a narrow Fe-K line. These 3 components were separated securely with the C3PO method. In
contrast, another Suzaku observation made in 2009 foundthis AGN in a fainter state with a 2–10 keV luminosityof 3.0 × 1042 erg s−1 (Noda et al. 2013), comparable tothe 3rd brightest data set in the present study. Like thepresent case, the primary continuum was then described
also by a single PL with Γ ∼ 1.7.
The issue is whether the PL with Γ = 1.7, identiﬁed in
the present data, as well as in 2009, is the SRPC or the
HGPC. Although it is neither as hard nor as strongly absorbed as the genuine HGPC (with Γ ∼ 1.5) co-existingwith the SRPC, we still regard it as the HGPC, because
it showed only rather weak intra-day variability.

213

1.4

6

7

2
1

5×Reflection0

Main hard PL

-1
2

2

3

4

5

1
0.8
0.6
0.4

No delay+2 days+4 days

0.2
0
-0.2
56380

56390

56400

56410 56420
Time (MJD)

56430

56440

56450

1.4

1.5

1.2

1
0.5
0

-0.5
56300

56400

56500

56600
MJD

56700

56800

B-band flux density (mJy)

2−10 keV flux
(10-11 erg s-1 cm-2)
Optical flux density
(mJy)

3

Epoch 1 2 3 4 5(a) Epoch 1

Scaled 2–10 keV flux

4

B-band flux density (mJy)

1.2

0.8
0.6
0.4
0.2
0
-0.2
0

Fig. 3. (top) Absorption-removed light curves of the 2–10 keV
ﬂux of the HGPC continuum (large open circles) from thepresent campaign, and that of the reﬂection component including the Fe-K line (small ﬁlled symbols). The latter ismultiplied by a factor or 5. (bottom) Light curves of the
B-band diﬀerential ﬂux density (see text) of the NGC 3516
nucleus, measured by the 5 telescopes (Noda et al. 2016).

3.3. Optical light curves
In the optical wavelength, NGC 3516 was also very faintduring the present campaign. This made the absolute
optical photometry diﬃcult, due to uncertainties in subtracting the host galaxy. Therefore, we decided to consider only relative changes of the optical ﬂux of the AGN,
employing the Diﬀerence Image Photometry technique(Tomany and Crotts 1996). Measurements with diﬀerent telescopes showed good mutual consistency.
Figure 3 (bottom) shows the B-band ﬂux densities of
the NGC 3516 nucleus, thus measured by the 5 telescopes, compared with the X-ray light curves in the top
panel. As stated above, the zero point of the optical ﬂux
density is not meaningful. Very fortunately, the source
varied signiﬁcantly in both frequencies, just when themonitoring was densest. Furthermore, the X-ray andoptical variations are very well correlated.
3.4. Optical vs. X-ray time lags
Our next task is to look for possible time delays between
the two bands. If the optical emission from the accretion disk varies and it causes changes in the seed photon
supply to the Comptonization process that presumably
produces the hard X-rays, we expect the X-ray variations
to lag behind those in the optical (Nandra et al. 2000).
If, in contrast, the optical intensity changes are mainly
due to reprocessing of the varying X-ray irradiation by
the accretion disk, the delay would be in the opposite(b)

1

No delay+2 days+4 days
0.5
1
1.5
2
2.5
2–10 keV flux (10-11 erg s-1 cm-2)

3

Fig. 4. (a) Enlarged light curves, from epochs 1 to 5, of the B band
(black) and the 2–10 keV HGPC (open red circles) signals. Filledgreen squares and open blue triangles are the X-ray data delayed by+2 and +4 days, respectively. (b) A scatter plot between the 2–10
keV HGPC ﬂux and the B-band ﬂux density, with time delays applied
to the X-ray data by 0 days (open red circles), 2 days (ﬁlled green
squares), and 4 days (open blue triangles). From Noda et al. (2016).

sense (Krolik et al. 1991).
For the above purpose, the X-ray and optical light
curves in the ﬁrst 5 epochs are expanded in Fig. 4a.
There, the X-ray data, shown in open red circles, appear to slightly precede those in the optical. If we purposely delay the X-ray data points by 2 days (ﬁlled green
squares), the correlation becomes better, but a delay by
4 days (open blue triangles) appears too much.
For a further study, Fig. 4b correlates the optical data,
against the X-ray signals purposely delayed by 0, 2, and4 days (in the same symbols as panel a). Again, delaying the X-ray data by 2 days maximizes the correlation. Through more sophisticated analyses, we havefound that the optical signal is delayed by 2.0 ± 0.6 days
behind X-rays (Noda et al. 2016). Therefore, the varying
optical emission must be due to X-ray reprocessing.
Another clue from Fig. 4 is the amplitudes of the Xray and optical signal increments. The unabsorbed 2–10
keV ﬂux varied by ΔFx ∼ 2 × 10−11 erg s−1 cm−2 ; so
did the B-band ﬂux density by ∼ 1 mJy. Multiplying
the latter by the B-band frequency of 7 × 1014 Hz, theoptical ﬂux increment (in νF ν sense) becomes ΔFopt ∼0.7 × 10−11 erg s−1 . Since this is a reasonable fraction ofΔFx , the reprocessing scenario is reinforced. In addition,
the lowest absolute optical intensity measured, thoughvery uncertain, is considered to be  ΔFopt . We cantherefore ignore the spontaneous optical emission fromthe disk, and consider only the X-ray reprocessing.

214(a) Lamppost corona + standard disk
B-bandX-ray

coronaHX
BHR(b) Hot accretion flow + truncated disk
X-ray

B-bandX-ray-luminous
regionBH~2 light days
coronaFig. 5. Disk/corona geometries in the “lamp-post” conﬁguration(panel a), and in the model invoking a truncated disk and a hot
accretion ﬂow inside (panel b). From Noda et al. (2016).

4. Discussion
During the present campaign, NGC 3516 was dim, and
the X-ray spectrum was dominated by the HGPC. We
detected a clear X-ray vs. optical correlation over ∼ 50days, together with an optical delay by τ = 2 days. Considering the previous lack of correlations in NGC 3516
(Maoz et al. 2002) when it was more luminous and presumably dominated by the SRPC, we suggest that the
HGPC in an AGN is correlated with its optical intensity,whereas the SRPC is not.
The measured τ means a light travel distance of rrp =5×1015 cm, or ∼ 500 Schwarzschild radii for MBH (§3.1).
Therefore, the X-ray to optical reprocessing matter was
located at a very large distance from the central BH.
These results clearly rule out the lamp-post scenario,
shown in Fig. 5a: we present qualitative reasons below,
leaving quantitative discussion to Noda et al. (2016).
The local temperature of a standard accretion disk
around a Schwarzschild BH is given, in units of eV, as

1/4
−3 

r
3 · 107 M  η 
T (r) ∼ 7
Rin
MBH
0.01
where r is the distance from the central BH. If the disk
continues down to r  rrp as postulated in the scenario,
these inner regions would produce strong optical emission, giving a large oﬀset to the optical intensity of thenucleus. This disagrees with our result, that the underlying optical ﬂux was  ΔFopt . In addition, the X-ray

reprocessing eﬃciency of the lamp-post model would increases steeply toward smaller r. This would make τmuch smaller than 3 days. In short, the lamp-post modelis unrealistic for NGC 3516, at least when it is dim.
A more successful alternative is given in Fig. 5b.
There, the disk is assumed to be truncated at r ∼ rrp ,inside which the accretion ﬂow turns into a hot corona,or Radiatively Ineﬃcient Accretion Flow (RIAF) whichis expected to appear at low η. By Comptonizing the
low-energy disk photons, the RIAF region will becomea hard X-ray source, but it would neither produce signiﬁcant optical emission, nor reprocess the varying hard
X-rays into visible light. Due to the lack of opticallythick material near ISCO, the X-ray spectra would bearno relativistic reﬂection signature. The value of Γ ∼ 1.7
is consistent with the RIAF model prediction. All theseproperties, expected for the geometry of Fig. 5b, agreewith the present observational results.
Finally, the above arguments are not speciﬁc to NGC
3516. In fact, several previous works that found relatively good X-ray vs. optical correlation reported optical lags by several days. These incude NGC 5548 (Suganuma et al. 2006; McHardy et al. 2014), NGC 3783
(Arévalo et al. 2009), and Mrk 79 (Breedt et al. 2009).
5. Conclusion
1. The optical emission from an AGN is correlated with
the X-ray HGPC, but not with the SRPC.
2. The lamp-post scenario, invoking an accretion diskthat extends close to the ISCO, is unrealistic, atleast when AGNs are dim.
3. The data favors an alternative scenario which invokes a truncated disk and a hot RIAF region.
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215

An X-ray spectral variability of fast disk winds in AGN
Kouichi Hagino,1 Hirokazu Odaka,2 Chris Done,1,3 Shin Watanabe,1,4 Tadayuki Takahashi1,41Institute of Space and Astronautical Science (ISAS), Japan Aerospace Exploration Agency (JAXA),
3-1-1 Yoshinodai, Chuo, Sagamihara, Kanagawa 252-5210, Japan
2
KIPAC, Stanford University, 452 Lomita Mall, Stanford, CA 94305, USA
3
Department of Physics, University of Durham, South Road, Durham DH1 3LE, UK
4
Department of Physics, University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan
E-mail(KH): hagino@astro.isas.jaxa.jp
Abstract

Recent X-ray observations of blue-shifted absorption lines revealed an existence of the extremely fastdisk winds with outﬂow velocities of ∼ 0.1–0.3c. Such fast outﬂows would have a large impact on the
coevolution of black holes and host galaxies since they are expected to carry a large amount of kinetic
energy. One of the common characteristics of these fast winds is a strong time variability of the absorption
feature. To investigate this variability, we have developed a new X-ray spectral model of the disk winds,
which is generated by 3-dimensional Monte Carlo radiation transfer simulations on the assumption of therealistic wind geometry. By applying our wind model to the multi-epoch X-ray data of an archetypal
wind source PDS 456, we ﬁnd the variability in the absorption line is explained by a change of thewind outﬂowing angle without any large variability in the mass outﬂow rate of the wind. This result
indicates that the fast disk winds are stable and that local hydrodynamic instabilities produce a large
time variability of the absorption line. Moreover, we also apply our wind model to the disk-line source
1H 0707−495. The characteristic Fe-K feature in this source is successfully reproduced our disk windmodel for all the observations.
Key words: black hole physics — radiative transfer — galaxies: individual (PDS 456, 1H 0707−495)

Introduction

Broad iron line Blueshiftedfrom the wind absorption
lines

Continuum
absorption    

Blueshifted absorption lines with velocities of 0.1–0.3c
from highly ionized iron ions have been discovered inmany local active galactic nuclei (AGN) by recent Xray observations (e.g., Chartas et al. 2002; Reeves et al.
2003; Pounds et al. 2003a; Pounds et al. 2003b). Theseabsorption lines indicate that heavy ions are outﬂowing with such extremely fast velocities from the vicinity of supermassive black holes at the center of galaxies.
These outﬂows are often referred to as ultra-fast outﬂows(UFOs) (Tombesi et al. 2010). Since they are expected tocarry large amounts of kinetic energy to the host galaxies, they are thought to play an important role in theco-evolution of the black holes and the galaxies (King &Pounds 2015).
In spite of its importance, physical properties of UFOs
are still unclear. Particularly, the origin of spectral
variation of the blueshifted absorption lines in UFOs
is unknown. As an example, spectra of PDS 456 observed in 2007 and 2013 by Suzaku are shown in Fig. 1.
These spectra clearly show spectral changes in both absorption lines and continuum absorption, which are a

PDS 456: 2007-02-24

1.

PDS 456: 2013-03-08 

Fig. 1. Observed spectra of an archetypal ultra-fast outﬂow in
PDS 456.

216

57­59 deg
Normalized ﬂux

Normalized ﬂux

45­47 deg
2

Total

1

Direct

0.4
0.3
0.2

Reprocess

0.1
0.04
0.03
0.02
0.01

2
1

0.4
0.3
0.2

Total
Direct0.1

Reprocess

0.04
0.03
0.02
6

8

10
Energy (keV)

12

14

0.01

6

8

10
Energy (keV)

12

14

Fig. 2. Simulated spectra with diﬀerent angle (Hagino et al. 2016).

common interesting characteristic in UFOs. Since theblueshifted absorption lines are usually interpreted as
K-shell transitions of He-like and/or H-like iron ions,
outﬂowing materials should be highly ionized with an
ionization parameter of ξ ≡ L/nR2 ∼ 103–4 . On theother hand, such highly ionized material is not able toexplain the strong continuum absorption as shown in
Fig. 1, which requires relatively low ionized materials
with ξ < 102 . Thus, these materials are often thought∼
to be cool clumps embedded in the hot highly ionized
outﬂow. These clumps are naturally expected in hydrodynamic simulations of the accretion disk winds due
to ionization instability and/or hydrodynamic instability (Krolik, McKee & Tarter 1981; Takeuchi, Ohsuga,
Mineshige 2014). While occultation by the cool clumps
can easily explain the spectral variability in continuum
absorption (Matzeu et al. 2016), there is no clear picture for explaining the spectral variation in blueshifted
absorption lines produced by the highly ionized outﬂow.
2. Monte Carlo radiation transfer simulations
We construct a new spectral model of the UFO in orderto understand the spectral variability of the highly ionized absorption lines (Hagino et al. 2015). In this model,
we assume the UV-line driving mechanism. This is thebest candidate of acceleration mechanism of UFOs because of its good eﬃciency in AGN environment. In the
UV-line driving mechanism, moderately ionized materials are accelerated by radiation pressure due to boundbound transitions with UV photons. Since cross sectionsof bound-bound transitions are 3–4 orders of magnitudelarger than that of Thomson scattering, this mechanism
can eﬃciently accelerate and launch the fast disk wind
if materials are appropriately ionized. Moreover, since
AGN radiation typically peaks at the UV band, which

moderately ionizes materials and causes radiation pressure via bound-bound transitions, AGN is suitable for
the UV-line driving.
We assume a biconical geometry, which was developed
for accretion disk winds in cataclysmic variables (Knigge,
Woods & Drew 1995). Based on the hydrodynamic simulations of the UV-line driven disk winds, a covering fraction of Ω/4π = 0.15 is adopted (Proga, Stone & Kallman
2000; Nomura et al. 2016). A maximum velocity of thedisk wind is determined by the observed absorption lines,
and its distribution in the disk wind is calculated following the CAK velocity law (Castor, Abbott & Klein 1975).
A density distribution is derived from the velocity distribution and conservation of mass. Then, an ionization
structure of the wind is calculated 1-dimensionally along
stream lines by utilizing XSTAR (Kallman & Bautista2001), which outputs ion abundances and electron temperatures in the wind. The radiation transfer calculationis performed in these density, velocity, ionization distributions.
We adopt Monte Carlo method for the radiation transfer calculation in the UV-line driven disk wind because its non-spherical geometry makes it very diﬃcult
to solve with analytical methods. The Monte Carloradiation transfer simulation is performed by utilizing
a simulation framework called MONACO (Odaka et
al. 2011). MONACO is a general-purpose framework,
which have been applied for synthesizing X-ray spectrafrom many kinds of astrophysical objects (e.g., Odaka
et al. 2014; Odaka et al. 2016; Furui et al. 2016; Hagino
et al. 2015; Hagino et al. 2016). In this code, all important physical processes in photoionized plasma are
implemented, namely photoionization, recombination,
photoexcitation, Compton scattering (Watanabe et al.
2006). In addition to these, special relativistic eﬀects,
which are essential in UFOs, are also considered.

217

fast, thinslow, dense

θ=46 deg

θ=58 degθ

X-ray

accretion disk

sourceFig. 3. A schematic view of angle dependences of the absorption
lines.

X-ray spectra calculated by our simulations are shownin Fig. 2. As shown in the left panel, blueshifted absorption lines and a broad emission line similar to theobservations are reproduced. More interestingly, when
the wind is observed from a larger viewing angle, theabsorption lines become deeper and broader as shown in
the right panel. Fig. 3 schematically show how the absorption line features depend on the observer’s viewing
angle. At a smaller viewing angle, only a thin and fastpart of the wind is observed. On the other hand, at alarger angle, a slow and dense part is observed as well as
the fast and thin part, and it makes the absorption lines
broader and deeper.
3. Application to an archetypal wind source PDS 456
We apply our spectral model to observational data ofan archetypal disk wind source PDS 456. This source isobserved by Suzaku for 5 times between 2007 and 2013,and most of these data clearly show blueshifted absorption lines (Reeves et al. 2009; Goﬀord et al. 2014). Also,
as already shown in Fig. 1, the absorption line featuresstrongly vary. It means that this source is one of thebest target to study the origin of the spectral variationin the highly ionized absorption lines.
All the Suzaku spectra of PDS 456 are successfully
reproduced by our wind model with a mass outﬂow
rate of Ṁwind /ṀEdd = 0.13 and a terminal velocity ofvwind ∼ 0.3c. As examples, we show the spectra andbest-ﬁt models for observations in 2007-02-24 and 201303-08 in Fig. 4. Between the observations, we change
only the viewing angle and the terminal velocity withoutchanging the global parameters such as the mass outﬂow
rate.
A change of the viewing angle, which successfully explains the strong spectral variability of absorption lines,
possibly relate to a ﬂapping or/and an inhomogeneity ofthe disk winds. They are naturally expected in hydrodynamic simulations of the winds due to hydrodynamic
instability (Proga, Stone & Kallman 2000; Krolik, McKee & Tarter 1981; Takeuchi, Ohsuga, Mineshige 2014).
Thus, our result indicates that the strong variability of
the highly ionized absorption lines in UFOs could originate from a local instability of the wind, not from a
global change of the wind.
4. Application to a disk-line source 1H 0707−495
As a next target, we choose a narrow line Seyfert 1 galaxy
1H 0707−495, which is famous for its extremely smeared
disk-line feature (Fabian et al. 2004). However, this diskline interpretation for a characteristic iron-K spectralfeature in this source requires extreme conditions. Since
the iron-K line is extremely smeared, a black hole spinmust be close to maximum. Moreover, an X-ray emitting
region (corona) must be located very close to the event
horizon to explain its strong reﬂection component. Also,a large iron abundance of 7–20 is required by its strong
iron line.
Alternatively, we interpret the iron-K feature as anabsorption feature produced by the UFO. This interpretation is very natural because the spectral features in
1H 0707−495 is very similar to those in PDS 456. Also,
this interpretation is supported by optical data of thissource. It requires super Eddington accretion (Done &Jin 2016), and in such a situation, the accretion disk
must not be a standard thin disk, which is assumed in
the disk-line interpretation.
We apply our disk wind model to all 15 spectra of
1H 0707−495 observed by XMM-Newton and Suzaku,
and successfully reproduce the structure at iron-K bandby changing only the viewing angle (Fig. 5). It meansthat the spectra of this source can be explained by the
UFO with Ṁwind /ṀEdd = 0.2 and vwind = 0.2c. In addition to these data, we also compare an extrapolation of
the best-ﬁt wind model for Obs15 (right panel of Fig. 5)
with the NuSTAR data (Kara et al. 2015). We ﬁnd thatthe higher energy data is also explained by our disk windmodel. This strongly supports our disk wind interpretation for this source.
5. Conclusions
We have constructed a new spectral model by calculating the radiation transfer in a realistic wind geometry
based on the UV-line driving mechanism. We ﬁrst applied our disk wind model to an archetypal wind source
PDS 456, and found that the strong spectral variation inthis source can be explained by a change of the viewing
angle without changing a mass outﬂow rate. It indicatesthat the spectral variability is due to a local instability
or inhomogeneity of the wind. As the next target, weapplied to the “disk-line” source 1H 0707−495. Strong        

218




Δθ=3°
χ2ν=127.9/107



2007



Δθ=10°
χ2ν=88.5/89



2013c






















  



Fig. 4. Suzaku spectra of PDS 456 compared with our spectral model. Δθ indicates a diﬀerence between the observer’s line of sight and an
inner stream line of the wind (Hagino et al. 2015).

Δθ=19.1°
χ2ν=52.5/55










Δθ=25.9°χ2ν=43.1/51
































Fig. 5. XMM-Newton spectra of 1H 0707−495 compared with our spectral model. Δθ indicates a diﬀerence between our line of sight andthe inner stream line of the wind (Hagino et al. 2016).

Fe-K features in all the spectra of this source observed
by XMM-Newton/Suzaku are successfully reproduced byour disk wind model. Moreover, higher energy spectra
obtained by NuSTAR are also explained by our disk windmodel.
References
Castor J. I., et al. 1975, ApJ 195, 157
Chartas G., et al. 2002, ApJ 579, 169
Done C., et al. 2016, MNRAS 460, 1716
Fabian A. C., et al. 2004, MNRAS 353, 1071
Furui S., et al. 2016, ApJ 818, 164
Goﬀord J., et al. 2014 ApJ, 784, 77
Hagino K., et al. 2015, MNRAS 446, 663
Hagino K., et al. 2016, MNRAS 461, 3954
Kallman T., et al. 2001, ApJS 133, 221
Kara E., et al. 2015, MNRAS 449, 234
King A., et al. 2015, ARAA 53, 115
Knigge C., et al. 1995, MNRAS 273, 225
Krolik J. H., et al. 1981, ApJ 249, 422
Matzeu G. A., et al. 2016, MNRAS 458, 1311

Nomura M., et al. 2016, PASJ 68, 16
Odaka H., et al. 2011, ApJ 740, 103
Odaka H., et al. 2014, ApJ 780, 38
Odaka H., et al. 2016, MNRAS 462, 2366
Pounds K. A., et al. 2003, MNRAS 345, 705
Pounds K. A., et al. 2003, MNRAS 346, 1025
Proga D., et al. 2000, ApJ 543, 686
Reeves J. N., et al. 2003, ApJ 593, L65
Reeves J. N., et al. 2009, ApJ 701, 493
Takeuchi S., et al. 2014, PASJ 66, 48
Tombesi F., et al. 2010, A&A 521, A57
Watanabe S., et al. 2006, ApJ 651, 421

219

220

221

Revisiting the cosmological evolution of X-ray selected blazars
with the Swift/BAT 70-month all sky survey data
Koyo Toda,1 Yoshiyuki Inoue,2 Yasuyuki Tanaka,1 and Yasushi Fukazawa1
1

Department of Physical Science, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526
2
Department of High Energy Astrophysics, Institute of Space and Astronautical Science,
Japan Aerospace Exploration Agency, 3-1-1 Yoshinodai, Sagamihara, Kanagawa 229-8510
E-mail(KT): toda@astro.hiroshima-u.ac.jp
Abstract
Blazars constitute the most extreme class of active galactic nuclei. Recent multi-wavelength observations allow us to understand the emission processes in blazars. In contrast, the cosmic history of blazars isstill veiled in mystery. The keys to understand their evolution is the cumulative source count distribution
and the luminosity function with which we are able to obtain the trend of evolution and the peak of blazar
formation epoch in the universe. In this talk, we report our study on the cosmological evolution of blazars
utilizing the 70-month data from the Swift/Burst Alert Telescope (BAT) survey. Our sample comprises 41
ﬂat-spectrum radio quasars (FSRQs) and 27 BL Lacertae (BL Lac) objects. Our sample size is a factor of∼ 2 bigger than that of the previous study with the 22-month Swift/BAT survey data (Ajello et al. 2009).
Although Ajello et al. 2009 reported the peak of the FSRQs density at z ∼ 4, we suggest that the FSRQs
density has a peak z < 4; that is, the redshift evolution of FSRQs is not stronger than the previous study.
Furthermore, we will discuss prospects for blazar detectability by future X-ray survey missions such as
eROSITA satellite.
Key words: AGN:Blazar — Luminosity function

1.

Introduction

Blazars constitute the most extreme class of active galactic nuclei (AGNs). Blazars have a relativistic jet that ispointing toward the line of sight and emit with radiationfrom radio waves to gamma-rays. Astronomers currently
recognize two types of blazars. Those are known as FSRQs and BL Lac objects. FSRQs have high luminosityand high redshift, and furthermore, its Equivalent widths(EW) of optical lines is greater than 5Å. In contrast, BL
lac objects have low Luminosity and low redshift, andfurthermore, its EW of optical lines is less than 5Å.
In previous study, the density of the gamma-ray
blazars peaks at z ∼ 1 − 2 in the history of the universe (Ajello et al. 2012). However, the density of X-ray
blazars peak are reported to be peaked at z ∼ 4 (Ajelloet al. 2009).
In the previous study, the number of X-ray blazar sample, which comprises 26 FSRQ samples from Swift/BAT
22-month catalog, is less than that of gamma-ray blazar
sample, which comprises 168 FSRQ from the Fermi ﬁrst
year catalogue. Then, it is important to derive number
density of FSRQs utilizing more X-ray blazars.

Therefore, the purpose of this study was to explore thecosmological evolution of blazars utilizing the 70-month
data from the Swift/BAT survey which is the newest
hard X-ray all sky survey map.
2. Sample Selection
In this work, we used the Swift BAT 70-Month Hard
X-ray Survey data (Baumgartner et al. 2013). This survey comprises 1211 X-ray objects which was detectedby Swift/BAT within 70 months. In order to securely
identify and classify BAT sources as blazars, we relied
mainly on the blazar catalog (BZCAT) (Massaro et al.
2009) which contains only bona ﬁde blazars. As a result,
our blazar sample comprises 41 FSRQs and 27 BL Lac
objects and the size of our sample is a factor of ∼ 2 bigger than that of the previous study with the 22-month
Swift/BAT survey data (Ajello et al. 2009).
3. Result3.1. logN-logS distribution
We derived logN-logS distribution. This distribution is
generally expressed as N (> S) = AS −β ,where S is a222

source ﬂux. If the source population is uniformly distributed, a β =1.5 is expected.
The logN-logS distributions we derived is shown in Figure 1 and Table 1. This logN-logS distribution indicatesthat FSRQs is positive evolution and BL Lac obejacts isnegative evolution.

3.2.2. Number density of FSRQ
Based on 2powMPDE and 1powMPLE model, we can seethat the peak of the FSRQ density at z < 4. This suggests that the redshift evolution of FSRQs is not stronger
than the previous study which used 22month Swift/BAT
data.

Fig. 1. LogN-LogS distribution of FSRQs (left) and BL Lacs (right).
The red line shows the result of previous study (Ajello et al. 2009)and black line shows the result of this study.

Table 1. The expected number of blazar detected by eROSITA satellite

Sample
FSRQs
BL Lac objects

βa1.87 ± 0.07
1.32 ± 0.13bβ22month
2.077 ± 0.269
1.694 ± 0.316

Note.
a Best-ﬁt exponent of the logN―logS distribution (e.g., N (> S) =AS −β )b Best-ﬁt exponent of the previous study (Ajello et al. 2009)

Fig. 2. Number density of FSRQs. The blue line are derived fromthe [15―55keV] LF (Ajello et al. 2009), the black line are derived
from the gamma-ray LF obtained by Fermi/LAT (Ajello et al.
2012) and the red line are derived from the X-ray LF obtained bythis work.

4. Blazar detectability by eROSITA satellite
From our logN-logS, we can estimate the expected number of blazers detected by eROSITA satellite. The results
are given in Table 2
Table 2. The expected number of blazar detected by eROSITA

3.2. Luminosity function of FSRQs
The luminosity function (LF) is deﬁned as the number
objects per unit comoving volume and per unit luminosity interval. The simple scenarios of evolution are pureluminosity evolution (PLE) and pure density evolution(PDE). We modeled the LF with eight widely used functions and applied the maximum likelihood method andKolmogorov-Smirnov (KS) test to derive the LF of FSRQs.
3.2.1. KS-test
In this work, 3 LF-models was not rejected:1powPDE
model, 2powMPDE model and 2powMPLE model.
These LF-models are listed below.

Observationtime
4 yearsurveySensitivity(erg/s/cm2 )2 × 10−13@2 − 10keV

FSRQs

BL LacObjects

13900

1900

Note. We assumed that logN-logS slope does not change in low
ﬂux range.

5. Conclusion
We suggest that the evolution of blazars is not stronger
than the previous study (Ajello et al. 2009)) and the peakof the FSRQs density at z < 4 . We will estimate thecontribution of blazars to the cosmic X-ray backgroundoutside the energy range band of this survey using LF.
References
Ajello, M. et al. 2009, Apj, 699, 603
Ajello, M. et al. 2012, Apj, 751, 108
Baumgartner, W. H. et al. 2013, ApJS, 207, 19B
Massaro, E. et al. 2009, A&A, 465, 691

223

MAXI and AKARI investigation of nearby active galactic nuclei
Naoki Isobe,1 Taiki Kawamuro,2 Shinki Oyabu,3 Takao Nakagawa,4 Shunsuke Baba,4,5
Kenichi Yano,4,5 Yoshihiro Ueda,2 and Yoshiki Toba6,7
1

Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro, Tokyo 152-8551, Japan
Department of Astronomy, Kyoto University, Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan
3
Graduate School of Science, Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602, Japan
4
Institute of Space and Astronautical Science (ISAS), Japan Aerospace Exploration Agency (JAXA),
3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 252-5210, Japan
5
Department of Physics, The University of Tokyo, 7-3-1, Hongo, Bunkyo-ku, Tokyo, 133-0033, Japan
6
Research Center for Space and Cosmic Evolution, Ehime University,
2-5 Bunkyo-cho, Matsuyama, Ehime 790-8577, Japan
7
Academia Sinica Institute of Astronomy and Astrophysics, PO Box 23-141, Taipei 10617, Taiwan
E-mail(NI): n-isobe@hp.phys.titech.ac.jp
2

Abstract
X-ray and infrared properties of nearby active galactic nuclei were studied with MAXI and the Japanese
infrared observatory AKARI, respectively. Among the 100 active galactic nuclei (except for blazars)
tabulated in the second MAXI catalog, 69 ones were found to have an infrared counterpart in the AKARI
point source catalog. The X-ray luminosity of these objects was corrected for absorption by referring to the
X-ray hardness ratio measured with MAXI. The absorption-corrected X-ray luminosity of Compton-thin
objects was found to be linearly correlated to their observed IR luminosity. This kind of correlation is
widely interpreted by the so-called clumpy torus model. In contrast, NGC 1365, only one Compton-thickobject in the MAXI catalog, was found to deviate from the correlation, toward a lower X-ray luminosity
by nearly an order of magnitude. It was veriﬁed that Compton-thick objects are eﬀectively picked up by
using the relation between the X-ray spectral hardness and X-ray-to-infrared color.
Key words: galaxies: active — galaxies: Seyfert — X-rays: galaxies — infrared: galaxies

1.

Introduction

Based on the uniﬁed picture (e.g., Antonucci 1993), it
is widely believed that the central engine of the active
galactic nuclei is a supermassive black hole accompanied
by a mass-accreting disk, and they are surrounded by aparsec-scale dust torus. However, our knowledge on theirdetailed spatial conﬁguration and physical condition still
remains far from satisfactory. Especially, the astrophysical importance of the Compton-thick objects, deﬁned asheavily obscured objects with a hydrogen column density of NH > 1.5 × 1024 cm−2 , has been widely recognized (e.g., the origin of the X-ray background; Ueda etal. 2014). Since these obscured sources are elusive in the
optical band due to signiﬁcant dust extinction, a useful
identiﬁcation technique is urgently requested.
It is widely thought that X-ray emission from active
galactic nuclei originates in the vicinity of the central
engine. These nuclear photons are absorbed in the dust
torus, and are re-emitted into the infrared (IR) band.
Therefore, a combination of X-ray and IR observations

provides us with a useful probe for the central region of
the active galactic nuclei. In addition, the X-ray and IR
color information is proposed to eﬀectively pick up theheavily obscured sources (Terashima et al. 2015).
The present study utilizes the second release of the
all-sky X-ray source catalog (hereafter the 2MAXI catalog; Hiroi et al. 2013) and the AKARI all-sky survey
point source catalog (Ishihara et al. 2010; Yamamura etal. 2012) to systematically diagnose nearby active galactic nuclei (Isobe et al. 2016). The sensitivity of the
2MAXI catalog is highest among the available all-sky
X-ray surveys. The IR spectrum is thought to be dominated by the torus emission in the 10–20 μm range, where
the AKARI catalog has a high sensitivity. These makethe combination of the 2MAXI and AKARI catalogs theideal tool for investigation of active galactic nuclei.
2. Results
The 2MAXI catalog lists 95 Seyfert galaxies and 5
quasars, located at the redshift of z = 0.002–0.15.

224

Fig. 1. (right) Relation between the absorption-corrected 4 – 10 keV X-ray luminosity, LH , and the 9 μm infrared luminosity, L9 . The ﬁlled
blue circles and open red diamonds indicate the type-1 and type-2 Seyfert galaxies, respectively. The Compton-thick object, NGC 1365, isplotted with the open triangle. The average of the X-ray to IR luminosity ratio in the logarithmic space for the type-1 objects is shown withthe solid line, while the 1σ and 2σ ranges are indicated by the dashed and dash-dotted lines, respectively. (left) X-ray hardness ratio, HR,plotted against the absorption-inclusive 3–4 keV X-ray to 18 μm IR luminosity ratio, LS /L18 . The model track for some representativevalues of the scattered X-ray fraction are drawn with the dotted lines as a function of the column density, NH , (the dashed line).

They are classiﬁed into 73 type-1 and 27 type-2 objects. Among them, 69 Seyfert galaxies, including 48
type-1 and 21 type-2 Seyfert galaxies, were identiﬁed as
an AKARI source.
In the left panel of Figure 1, the 4–10 keV X-ray luminosity (LH ) is plotted against the 9 μm IR luminosity(L9 ) for the Seyfert galaxies detected with both MAXI
and AKARI. Here, the X-ray luminosity was corrected
for absorption by utilizing the hydrogen column density
which is inferred from the X-ray hardness ratio, HR, tabulated in the 2MAXI catalog (Ueda et al. 2011), while
no correction was applied to the IR luminosity. Figure 1
strongly supports that the X-ray luminosity linearly correlates to the IR ones, regardless of the classiﬁcation of
the active galactic nuclei, at least for the Compton-thin
objects.
A similar X-ray-to-IR correlation was commonly foundin recent studies of nearby active galactic nuclei (Gandhi
et al. 2009; Matsuta et al. 2012). A simple torus modelwith a smooth and homogeneous dust distribution fails
to interpret this relation, because dust extinction to the
IR luminosity is thought to become not negligible for
type-2 objects with a high NH value. Therefore, the socalled clumpy torus geometry (e.g., Krolik & Begelman1988) is widely proposed.
In the 2MAXI catalog, one Compton-thick object,
NGC 1365, is listed. It is important to note that thissource signiﬁcantly deviates from the LH –L9 correlation
for the Compton-thin sources, as indicated by the red triangles in the left panel of Figure 1. This is because theabsorption correction by the MAXI hardness ratio is only
applicable to Compton-thin sources, of which the X-ray
emission is dominated by the direct/absorved component

from their nucleus. The method inevitably underestimate the intrinsic luminosity of Compton-thick objects,
since their dominant spectral component is though to be
reﬂected and/or scattered emission (Ueda et al. 2014).
This implies that the relation of the X-ray hardness
to the X-ray-to-IR color is utilized as an eﬀective tool
to ﬁnd Compton-thick active galactic nuclei (Terashima
et al. 2015). In the right panel of Figure 1, the HR
value derived with MAXI is plotted against the ratio of
the 3–4 keV luminosity to the 18 μm IR one, LS /L18 ,for the 2MAXI-AKARI sources. With the direct andscattered/reﬂected components taken into account, themodel spectrum is possible to qualitatively reproduce theproperties of both the Compton-thin and -thick sourceson this diagram, as shown with the dotted lines in
the Figure 1 right panel. Especially, the ﬁgure conﬁrms the ability of this kind of diagrams for discriminating Compton-thick objects from Compton-thin ones.
This method will become a powerful tool to search forCompton-thick sources in future X-ray and IR surveys.
References
Antonucci, R. 1993, ARA&A, 31, 473
Hiroi, K. et al. 2013, ApJS, 207, 36
Gandhi, P. et al. 2009, A&A, 502, 457
Isobe, N. et al., 2016, PASJ, 68, 98
Krolik, J. H., & Begelman, M. C. 1988, ApJ, 329, 702
Matsuta, K. et al. 2012, ApJ, 753, 104
Terashima, Y. et al. 2015, ApJ, 814, 11
Ueda, Y., et al. 2011, PASJ, 63, S937
Ueda, Y. et al. 2014, ApJ, 786, 104
Ishihara, D. et al., 2010, A&A, 502, 457
Yamamura, I. et al. 2012, PKAS, 27, 107

225

Kanata optical monitoring of radio galaxy NGC 1275
Kensei Shiki,1 Y. Fukazawa,1 Y. Tanaka,1 R. Itoh,1 2 K. Takaki,1 M. Kawabata,1
T. Nakaoka,1 T. Abe,1 N. Kawahara,1 J. Hirochi,11Department of Physical Science, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526
2

Department of Physics, Tokyo Institute of Technology, Ookayama, Meguro-ku, Tokyo 152-8551
E-mail(FS): shiki@astro.hiroshima-u.ac.jp
Abstract

NGC 1275 is a radio galaxy, located at the center of Perseus cluster. The multi-wavelength spectrum
from radio to gamma-ray of NGC 1275 is similar to that of blazars, and it shows a variable gamma-ray
activity. NGC 1275 showed broad emission lines in optical band and a Fe-K line in X-ray band which aresimilar to Seyfert galaxies. Therefore emission of accretion disk dominates in optical and X-ray bands.
In previous works (e.g., Yamazaki et al. 2013), there was no evidence about time variability in the
optical bands from 2010 to 2011, and X-ray bands from 2006 to 2011. On the other hand, Aleksic etal. (2014) showed a variability in gamma-ray and optical bands by the data of MAGIC between October
2009 and February 2010, Fermi-LAT and KVA in between and August 2010 and February 2011. In theirpaper, optical variability was also reported.
Fukazawa et al. 2016 (astro-ph) has reported graduallybrightening of NGC 1275 in X-ray correlated with gamma-ray activity from 2013 to 2014. Based on theseresults, we expect a jet origin brightening in optical band during the same time period. We performed
optical and near-infrared photometry using HONIR installed on the 1.5 m Kanata telescope since February
2015. Furthermore, we analyzed X-ray archival public data of XRT onboard the Swift satellite, and also
investigated time variability correlation between X-ray and optical bands.
Key words: AGN Jet, Radio Galaxy, Kanata Telescope, Swift, Fermi

1. NGC 1275
NGC 1275 is a radio galaxy located at the center of Perseus cluster, and object position is RA =03h:19m:36.21s, Dec = +41:29 :57.9 . As shown in
Fig. 1, the emission from NGC 1275 has been observedin the broad-band energy range from radio to gammarays. NGC 1275 showed emission lines in optical band
and X-ray band, suggesting that the emission associated with the accretion ﬂow exists. On the other hands,
there are strong variabilities in gamma-ray band. Recently, MAGIC and VERITAS (both of them are Atmospheric Cherenkov Telescope) reported the detectionof very-high-energy gamma-ray activity from NGC 1275
during the same epoch in Oct. 2016. Therefore, the jet
must contribute to the emission of NGC 1275 in the radio and gamma-ray band. But the radiation mechanism
and the structure of jets have not been understood yet.
2. Data Reduction
2.1. Swift satellite and XRTThe Swift is a multi-wavelength observation satellite
launched on November 20, 2004. Swift has three instruments; Burst Alert Telescope (BAT), X-ray Telescope

Fig. 1: The spectral energy distribution of NGC 1275
(S. Yamazaki et al 2013)(XRT), and UV/Optical Telescope (UVOT). The XRT
focuses on X-ray with 23.6 x 23.6 arcmin FOV and 0.2-10
keV energy range.
2.2. X-Ray Data Analysis
We used public data of Swift XRT obtained from July
2007 to November 2016. The bright emission from the
Perseus cluster was also observed, and thus we extracted
the XRT spectrum within 0.3 arcmin of NGC 1275, and226

1.0e-06
8.0e-076.0e-074.0e-072.0e-07

Kanata telescope and HONIR

Swift XRT

5.5e-11FLUX [erg/cm2/s]

Kanata is 1.5m Ritchy-Cretien telescope at HigashiHiroshima Observatory. It has 1 Cassegrainfocus and 2
Nasmythfocuses. Hiroshima Optical and Near-Infrared
camera (HONIR), installed on the Cassegrainfocus, cantake image in the optical and near-infrared (NIR) band
simultaneously. Observed targets are focused especiallyon variable objects such as blazars, supernovae, andgamma-ray bursts.

5.0e-11
4.5e-11
4.0e-113.5e-113.0e-112.5e-112.0e-111.5e-111.0e-11
54000

54500

55000

55500

56000

56500

57000

57500

58000

Fig. 2: Fermi and Swift whole period light-curve

Optical observation and Analysis

We took data on 45 nights in February 2015 - November 2016. Optical data were obtained in the V and R
bands, and NIR data in the J bands. The host galaxy
and extended nebula are very bright, and thus it is hardto subtract the emission from these components. Table
1 shows the number of data obtained with HONIR.

2.0e-06[photn/cm2/s]

2.4.

Fermi 1week bin

Fermi 1day bin
Fermi 1week bin
1.0e-06

6.0e-11[erg/cm2/s]

2.3.

1.2e-06

FLUX [photn/cm2/s]

performed the spectral analysis of AGN. The background
spectrum was obtained at the position of 60–65 arcsec from NGC 1275, and then subtracted from NGC
1275 spectrum. We ﬁtted all the XRT spectra withwabs*pegpwrlw in the XSPEC model.

4.0e-11

V
42
150

R
36
141

J
42
151

Ks
36
141

0.2

R
V

0
0.4

normed mag

year
2015
2016

normed mag

0.4

Table 1: Kanata Telescope photometry data list

Swift XRT

J

0.2
0

57300

57400

57500

57600

57700

57800

TIME [MJD]

3.

Results Of X-ray And Optical Data

Fig. 2 shows light curves, based on all Fermi achival
gamma-ray data and Swift X-ray data. There are strong
variabilities in gamma-ray band. We can see signiﬁcantbrightness increase around MJD 5550 and 57000 in both
gamma-ray and X-ray. Fig. 3 shows multi-wavelength
light-curve from 57300 to 57800. In s short time-scale,no obvious synchronized feature among the bands can be
seen.

4.

Summary and Feature work

In this work, we analyzed X-ray data of Swift-XRT and
optical and NIR data of HONIR. We compared these results with gamma-ray archive light curve of Fermi-LAT.
Currently we do not see a signiﬁcant correlation between
optical and Gamma-ray for the short-term gamma-ray
variability. Because the optical ﬂux estimation is still
not optimized in this work, we are improving the optical
ﬂux estimation.

Fig. 3: Multi-wavelength light-curve from 57300 to
57800[MJD]. From the top of ﬁgure, Fermi gamma-ray,
Swift X-ray, Kanata optical and Kanata near infrared.
5. References
References
Aleksic, J., et al. March 6, 2014, A-2013-22951 arxiv
Yamazaki, S., et al. ”X-Ray and Optical Monitoring of a
Gamma-Ray-Emitting Radio Galaxy,NGC 1275” 25
April 2013
Fukazawa, Y., et al. ”X-ray variability of GeV gammaray emitting radio galaxy NGC 1275” 12 August 2016
The Astronomer’s Telegram #9689
The Astronomer’s Telegram #9690

227

X-Ray Spectral Variability of the Narrow-Line Seyfert 1 Galaxy
NGC 4051 Observed with Suzaku
Manami Seino,1 Hirofumi Noda,2 Katsuma Miyake,1 Kazuhiro Nakazawa1 and Kazuo Makishima3
1

The University of Tokyo, 2 Tohoku University, 3 RIKEN
E-mail(MS): seino@juno.phys.s.u-tokyo.ac.jp
Abstract

The narrow-line Seyfert 1 galaxy NGC 4051 was observed with Suzaku twice in 2008 November, withan interval of 2 weeks. It showed a high X-ray ﬂux of ∼ 2 × 10−11 erg cm−2 s−1 at 2–10 keV, which
varied by a factor of 2 in less than 500 s. To study the source variation on time scales of ∼ 500 s, we took
the spectral diﬀerence between bright and faint phases in one observation. This diﬀerence spectrum wasrepresented by a single power-law (PL) with photon index Γ ∼ 2.2. However, a simultaneous ﬁtting to
the time-averaged and diﬀerence spectra failed if only adding a reﬂection component to the Γ ∼ 2.2 PL.
By adding another harder PL as in Seyfert 1 galaxies (e.g. Noda et al. 2013), the ﬁt became successful.
Furthermore, the diﬀerence spectrum between the two observations showed a clear hard excess component
in addition to the Γ ∼ 2.2 PL. The result suggests that the new PL ﬂux also varied in 2 weeks, whichindicates a source size of 60Rg < R < 105 Rg .
Key words: galaxies: active — galaxies: individual (NGC 4051) — X-rays: galaxies

1. Introduction
The X-ray spectra of active galactic nuclei are considered
to consist of a power-law (PL) component with photon
index Γ ∼ 2, which varies on a time scale of ∼ ks, andits reﬂection component with Fe-K line which varies in ∼month (Fabian & Miniutti 2005). However, this common
view has not been fully tested against the most up-todate data. Actually, Noda et al. (2011; 2013; 2014)
revealed that more than one primary components exist
in Seyfert 1 galaxies (Sy1s).
In this work, we investigated if the same view can beapplied to narrow-line Seyfert 1 galaxies (NLS1s), whose
black holes (BHs) are less massive and Eddington ratios
are higher than those of Sy1s. NGC 4051 is a near-by(redshift z = 0.0024; Brinkmann et al. 1995) NLS1, with
BH mass of ∼ 1.7 × 106 M (Denney et al. 2009). Beingbright and highly variable, this is one of the best targets
for our study.

Fig. 1. Background-subtracted light curves of NGC 4051, obtainedwith XIS0+3 (2.0–10.0 keV in black) and HXD-PIN (15–45 keV
in gray), shown with a binning of 3 ks. The lower panel shows
the same XIS0+3 light curve in Obs.1 with 500 s binning, wherevertical lines deﬁne data subsets.

2. Observation & Analysis
NGC 4051 was observed with Suzaku twice in 2008
November, with an interval of 2 weeks. We call theobservations Obs.1 and 2 in chronological order (Fig.1)
with exposures of 275 and 78 ks, respectivery. In thepresent work, we utilized data from the X-ray Imaging
Spectrometer (XIS) and the Hard X-ray Detector with
Si PIN photo-diodes (HXD-PIN). XIS events were ex-

tracted from a circular region of 3 .3 radius centered on
the source, and background events were from an annularregion of the inner and outer radii of 5 .9 and 8 .4, respectively. The data of XIS0 and XIS3 (XIS0+3) using
front-illuminated CCDs were added. We subtracted the
HXD-PIN background provided by the HXD team, andmodeled the Cosmic X-ray background based on Boldt

228

Fig. 2. (a)Time-averaged (black) and 500 s diﬀerence (red) spectraof NGC 4051 from Obs.1 obtained by XIS0+3 (2–10 keV) and PIN(15–45 keV). They are divided by a Γ = 2.0 PL. (b)Time-averaged
spectra of Obs.1 (black) and Obs.2 (gray), and their diﬀerence(red).

(1987).
In Fig.2, we present time-averaged spectra of Obs.1
and 2, all divided by a prediction for a Γ = 2.0 PL. The
spectral shapes of Obs.1 and Obs.2 are alike, and theintensity was 1.4 times higher in Obs.1 than in Obs.2.
We divided the observation time of Obs.1 into 20–30 ks
subsets (as shown in the lower panel of Fig.1), and made
spectra of fainter and brighter periods than the average in each subset using the 500 s bin light curve. The
red spectrum in Fig.2a is an overall diﬀerence spectrum,
created by averaging 23 diﬀerence spectra from the individual subsets. It has a shape of Γ ∼ 2.2 PL, and lacks
the Fe-K line and the hard X-ray hump.
Then, we focus on the energy band above 2 keV

for model ﬁtting. We simultaneously ﬁtted the timeaveraged and 500 s diﬀerence spectra of Obs.2. The difference spectrum was represented by a variable model(VM): PL1 × abs.(6.8 keV) × abs.(7.1 keV), and thetime-averaged spectrum by the sum of VM and a stablemodel (SM). VM has the same ﬁtting parameters as the
VM except its normalization. We multiplied the wholemodel by galactic and ionized absorption.
First, we applied only a reﬂection component as SM,
whose abundance and inclination angle were ﬁxed at 1
solar and 60◦ , respectively. As a result, the diﬀerence
spectrum was well explained by the VM, but some residuals remained in the time-averaged spectrum. Hence,
we needed to add another component to SM. As thisthird component, we examined three candidates: another PL2 with absorption, partial covering absorption of
the variable PL1 , and relativistically smeared reﬂectionof PL1 . As a result of ﬁtting, independent PL2 model
best represented the spectra (χ2 /d.o.f = 272/243, partial covering: χ2 /d.o.f = 296/244, relativistic reﬂection:
χ2 /d.o.f = 292/244). This component was a hard PL ofΓ = 1.5 ± 0.3 with equivalent hydrogen column density
NH = (5 ± 4) × 1022 cm−2 (Fig.3).
Furthermore, in the diﬀerence spectrum between thetwo observations, a hard X-ray hump in addition to the
PL was seen (Fig.2b). Because there was no Fe K-line
in the diﬀerence spectra, this hard X-ray hump can besaid not to be the reﬂection component.
3. Discussion & Conclusion
The hard X-ray hump of the diﬀerence spectrum (Fig.2b)
suggests the third component also varied in 2 weeks.
From this time variability, we discuss the properties of
the third component. If the hard PL is the right model,
the time variability suggests that a size of the sourceis 60Rg < R < 105 Rg . If partial covering, further ﬁnetuning to explain both the spectral diﬀerence and time
variability is needed. If relativistic reﬂection, we need ageometry to explain the 500 s invariability as well as the
2 weeks variability. Therefore, taking into account ourmodel ﬁtting in previous section, we conclude the hard
PL emission is most reasonable.
References
Boldt, E. 1987, IAUS, 124, 611
Brinkmann, W. et al. 1995, A&AS, 109, 417
Denney, K. D. et al. 2009, ApJ, 702, 1353
Fabian & Miniutti 2005, arXiv:astro-ph/0507409
Miyake, K. et al. 2016, PASJ, 68, S28
Noda, H. et al. 2011, PASJ, 63, 449

Fig. 3. Simultaneous ﬁt to the time-averaged and 500 s diﬀerence
spectra with the model described in text.

Noda, H. et al. 2013, ApJ, 771, 100
Noda, H. et al. 2014, ApJ, 794, 2

229

The new primary X-ray spectral component of IC4329A
conﬁrmed with Suzaku and NuSTAR
Katsuma Miyake,1 Hirofumi Noda,2 Shin’ya Yamada,3 Kazuo Makishima4 and Kazuhiro Nakazawa1
1

University of Tokyo, 2 Tohoku University, 3 Tokyo Metropolitan University, 4 RIKEN
E-mail(KM): miyake@juno.phys.s.u-tokyo.ac.jp
Abstract

The bright and highly variable Seyfert I active galactic nucleus, IC4329A, was observed simultaneouslywith Suzaku and NuSTAR for about 3 days. Within the observation, the 2–10 keV ﬂux changed by afactor of 1.5 around the mean of ∼ 1.0 × 10−10 erg cm−2 s−1 . Using the model-independent spectral
decomposition method developed by Noda et al. (2011), the time-averaged Suzaku and NuSTAR spectra(3–70 keV) were consistently decomposed into a fast and a slowly varying components. While the former
can be modeled as a power-law (PL) with a photon index Γ = 2.1, the latter was interpreted as the sumof a harder primary PL (Γ = 1.4) and the reﬂection. By combining the Suzaku and NuSTAR data, thepresence of this new primary PL, ﬁrst revealed in the Suzaku data (Miyake et al. 2016), has been reinforced.
Key words: galaxies: active — galaxies: individual (IC4329A) — galaxies: Seyfert — X-rays: galaxies

43
2Count rate (cnt/s)

0

50

100

150 200
Time (ks)

250

300

0.1

Fig. 1. Background subtracted and dead-time corrected light curvesof IC4329A in 2–10 keV with FPMA (black) and XIS FI (gray).
Time 0 is 2012 August 13 UT 02:13:09.

ratio
0.05

2. Observation and Data Reduction
IC4329A is a bright and highly variable Seyfert I AGN(Perola et al. 1999). This source was observed in 2012
August 13 simultaneously with Suzaku and NuSTAR.
We had already analyzed this Suzaku data set, and conﬁrmed the Γ = 1.4 PL in the 2–45 keV spectra (Miyake
et al. 2016). In the present work, we try to strengthen
the result by incorporating the NuSTAR data. We analyzed publicly available XIS 0, XIS 3, and HXD-PIN
data of Suzaku, and FPMA and B data of NuSTAR. XIS
0 and 3 data were co-added and are referred to as XIS
FI. On-source events of the XIS and FPM were extracted
from a circular region of 3 radius, and background eventsfrom an annular region of the radii 4 .8–7 .8. From the
HXD-PIN data, we subtracted Non X-ray Backgroundestimated by Fukazawa et al. (2009) and Cosmic X-ray
Background (Boldt 1987).

0.02

1. Introduction
X-ray spectra of Seyfert I active galactic nuclei (AGNs)
were considered to consist of a power-law (PL) like primary component, and its reﬂection component accompanied by an Fe Kα line (Fabian & Miniutti 2005). However, it is in reality diﬃcult to observationally extract the
primary PLs, because diﬀerent components are mixed in
an observed spectrum (Cerruti et al. 2011). Noda etal. (2011, 2013a) developed a model-independent spectral decomposition method, the “Count-Count Correlation with Positive Oﬀset” (C3PO) method. Using it, wediscovered that X-ray spectra of many AGNs harbor thetwo distinct primary PLs (Noda et al. 2013b, 2014).

5

20
10
Energy (keV)

50

Fig. 2. Background-subtracted time-averated spectra of IC4329A
obtained with FPMA (black), and XIS FI and the HXD (gray).

3.2. C3PO method
We applied the C3PO method to NuSTAR data for theﬁrst time. The 2–3 keV band of XIS FI was selected
as the reference band, because it exhibited the highest variability. Figure 3 shows examples of Count-Count
Plots (CCPs). Since all CCPs in ﬁner bands from 3 to
70 keV showed good linearity like ﬁgure 3, we applied the
C3PO method to the present data. As shown in ﬁgure 4,
we successfully decomposed the time-averaged spectra ofIC4329A into fast (FS) and slowly variable spectra (SS).
The result of the decomposition was consistent between
the NuSTAR and Suzaku data. The FS has a PL shapewith Γ ∼ 2, while the SS has the Fe K line.
FPMA count rate (cnt/s)

1.2
3-5 keV0.6

15-50 keV

0.8
0.4
0.4

0

0.2

0.4

0
0.4
0.8
1.2
XIS 2-3 keV count rate (cnt/s)

0.8

1.2

Fig. 3. CCPs between the reference XIS 2–3 keV and a broad FPMA
band. The best-ﬁtting straight line is shown in gray.

The C3PO method performs a simultaneous ﬁtting using two model components denoted as model F
and model S. The FS is ﬁtted with model F, the SSwith model S, and the time-averaged spectrum withmodel F+model S. As discussed in Miyake et al. (2016),
the SS of the Suzaku data requires not only the reﬂection but also another PL. Therefore, we adopted a PL
as model F, and again the sum of reﬂection and anotherPL as model S. As a result of the simultaneous ﬁttingwith parameters common to the NuSTAR and Suzaku
data, the ﬁt was successful with χ2 /d. o. f. = 427/368.

0.1
0.01
10−3
10−4
2

χ

3. Data Analysis
3.1. Light curves and spectra
Figure 1 shows background subtracted and dead-timecorrected light curves in 2–10 keV obtained with NuSTAR and Suzaku, with net exposures of 162 and 102 ks,
respectively. The source varied by a factor of ∼ 1.5
in the observation, around a mean 2–10 keV ﬂux of∼ 1.0 × 10−10 erg cm−2 s−1 . In the present work, weutilize the NuSTAR data only over the period which issimultaneous with Suzaku; the net exposure is reduced
to 68 ks. Figure 2 shows background-subtracted timeaveraged spectra in the form of their ratios to a PL modelwith photon index Γ = 2. In addition to an approximately PL-like continuum with a cutoﬀ at ≥ 40 keV,they exhibit Fe Kα line and edge at 6.4 and 7.2 keV in
the rest frame, respectively.

keV 2 (Photons cm -2 s -1 keV -1)

230

0
−2
5

10

20

50

5

10

20

50

Energy ( (keV)

Fig. 4. The time-averaged spectra (black) shown with the FSs (green)
and SSs (red), ﬁtted with the sum of the fast variable PL (green),
the slowly variable PL (magenta), and the reﬂection (blue).

The PL of model F has Γ = 2.0, while model S consistsof another Γ = 1.5 PL, strongly-absorbed with a column
density NH = 3.7 × 1023 cm−2 , and reﬂection with theabundance ﬁxed at 1 solar and the solid angle of 1.0π.
4. Discussion and Conclusion
The spectra of IC4329A was successfully decomposedwith the C3PO method, and they were well reproducedby considering two primary PL radiation as other AGNs.
So far, the NuSTAR spectrum of IC4329A has been
considered to include a relativistically smeared Fe line
(e.g. Brenneman et al. 2014). When the relativistic
reﬂection was adopted instead of the slowly variable PL,
the ﬁt became unaccepted with χ2 /d. o. f. = 648/367.
Thus, the advantage of our modeling over the relativistic
Fe line interpretation has been very much strengthened
by incorporating the NuSTAR data, because these modelings gave ﬁt χ2 /d. o. f. of 200/168 and 205/169, respectively, when using only the Suzaku data. Therefore, our“two primary PLs” interpretation has been reinforced by
combining both the NuSTAR and Suzaku data.
References
Boldt, E. 1987, Observational Cosmology, 124, 611
Brenneman et al. 2014, ApJ, 788, 61
Cerruti, M. et al. 2011, A&A, 535, A113
Fabian, A. C., & Miniutti, G. 2005, arXiv:astroph/0507409
Fukazawa, Y. et al. 2009, PASJ, 61, 17
Miyake, K. et al. 2016, PASJ, 68, S28
Noda, H. et al. 2013, ApJ, 771, 100
Noda, H. et al. 2013b, PASJ, 65, 4Noda, H. et al. 2011, PASJ, 63, 449
Noda, H. et al. 2014, ApJ, 794, 2
Perola, G. C. et al. 1999, A&A, 351, 937

231

Variability Analysis of X-ray spectrum of Blazar OJ 287 from Suzaku/XIS
and Swift/XRT (2005–2016)
Siti Fatima1
1

Department of Astronomy, Institut Teknologi BandungE-mail(SF): sitifatima@students.itb.ac.id
Abstract

Blazar OJ 287 is known to exhibit periodic outburst from optical observations. A model of BinarySupermassive Black Holes (SMBHs) in the center of OJ 287 is used to explain the phenomenon of 12
years periodic optic outburst. X-ray observations also conﬁrmed outbursts of OJ 287 in the X-ray band.
X-ray spectra of OJ 287 from X-ray observation in 2005 to 2016 were analyzed to determine the variability
in X-ray. Spectral data are obtained from X-ray Imaging Spectrometers (XIS) from Suzaku mission in
observation years of 2007 and 2015. Also obtained from X-ray Telescope (XRT) mission from Swift inobservation period since May 2005 to March 2016. Energy interval of the X-ray spectrum is 0.4 – 10
keV for Suzaku/XIS and 0.3 – 7 keV for Swift/XRT. The data are divided into several non-outburst andoutburst data groups. In general power law model is the best model and it is also conﬁrmed that the
X-ray spectra of OJ 287 during that period exhibit ﬂuctuation in photon index value, Γ. Variability of thevalue of Γ indicates that there is evolution of the X-ray spectrum of OJ 287 during non-outbursts time andalso during transition between non-outburst and outburst phase and vice versa. Diﬀerent values of Γ arealso found during diﬀerent outburst phases. In general the derived values of Γ are in agreement with theinverse compton model (Γ ∼1 – 2) in contrast with the synchrotron model (Γ ∼2 – 3). We also found thatΓ increases as the ﬂux increases as opposed to that of synchrotron model. In addition, X-ray luminosity
of OJ 287 shows ﬂuctuation within ∼ 1 × 1011 ≤ L ≤ 1 × 1012 .

232

Synergetic Growth of the Rayleigh-Taylor and Richtmyer-Meshkov
Instabilities in the Relativistic Jet
Jin Matsumoto1
1

RIKEN
E-mail(JM): jin.matsumoto@riken.jp
Abstract
The relativistic jet is ubiquitous phenomenon among astrophysical systems consisting of a compact
object surrounded by an accretion disk, e.g., active galactic nuclei, microquasars, and the central engine of
gamma-ray bursts. The dynamics and stability of the relativistic jet is important in order to understandthe emission properties of the jet. The radial oscillating motion of the jet is naturally excited due tothe pressure mismatch between the jet and surrounding medium, that is, cocoon when the jet propagates
through an ambient medium. In the rest frame of the decelerating jet interface that expands radially, aneﬀective inertia force acts on the interface and is directed outward. Therefore the jet medium is on top of
the cocoon medium in the eﬀective gravity in this frame and the Rayleigh-Taylor instability is expected
to grow at the interface. We investigate the growth of the Rayleigh-Taylor instability at the interface ofthe jet using three-dimensional hydrodynamic simulations. The propagation of the relativistic jet thatis continuously injected from the boundary of the calculation domain into a uniform ambient medium issolved. We ﬁnd that the interface of the jet is deformed by a synergetic growth of the Rayleigh-Taylor andRichtmyer-Meshkov instabilities regardless of the launching condition, such as the speciﬁc enthalpy of thejet or the eﬀective inertia ratio between the jet and ambient medium. The material mixing between the
jet and external medium due to these instabilities causes the deceleration of the jet.

233

Multi-wavelength observation of blazar Mrk 421 in extreme X-ray ﬂare
Ryosuke Itoh1
1

Tokyo Institute of technology
E-mail(RI): itoh@hp.phys.titech.ac.jp
Abstract
Blazars are highly variable active galactic nuclei that can be detected for all wavelengths and thought
to have relativistic jets. Blazar Mrk 421 exhibited extreme X-ray ﬂares in 2010. We performed opticalphotopolarimetric follow-up observations using the Kanata telescope. In 2010, the variability in the Xray band was signiﬁcant, while the optical and ultraviolet (UV) ﬂux decreased gradually. Polarization
properties also exhibited unique variability in 2010, suggesting the presence of systematic component ofpolarization and magnetic ﬁeld alignment for the emergence of a new polarized emission region. The results
revealed diﬀerent behaviors in terms of spectral evolution and suggested diﬀerent variability mechanisms
between 2010 and 2011. In 2010, the radiation was likely the result of energy injection into the emitting
regions with an aligned magnetic ﬁeld. In contrast, in 2011 the superposition of diﬀerent emission regions
may have contributed to the low degree of observed polarization. It also implies that high-energy electron
which were not accelerated to ultra-relativistic velocities were injected in 2011.

234

235

Gravitational
Wave
Sources

236

237

Binary black hole remnants of ﬁrst stars for the gravitational wave source
Tomoya Kinugawa1
1

Institute for Cosmic Ray Research, Univercity of Tokyo
E-mail(FS): kinugawa@icrr.u-tokyo.ac.jp
Abstract

Using our population synthesis code, we found that the typical chirp mass of binary black holes (BHBHs) whose origin is the ﬁrst star (Pop III) is ∼ 30 M with the total mass of 60 M so that the inspiral
chirp signal as well as quasi normal mode (QNM) of the merging black hole are interesting targets of
LIGO,VIRGO and KAGRA. The detection rate of the coalescing Pop III BH-BHs is ∼ 180 events/yr(SF Rp /(10−2.5 M /yr/Mpc3 ))·([fb /(1 + fb )]/0.33) · Errsys in our standard model where SF Rp , fb andErrsys are the peak value of the Pop III star formation rate, the binary fraction and the systematic error
with Errsys = 1 for our standard model, respectively. Furthermore, We found that the chirp mass has apeak at 30 M in most of parameters and distribution functions. This result predicted the gravitational
wave events like GW150914 and LIGO paper saidʞ recently predicted BBH total masses agree astonishingly
well with GW150914 and can have suﬃciently long merger times to occur in the nearby universe (Kinugawa
et al. 2014) ʟ. Thus, there is a good chance to check indirectly the existence of Pop III massive stars bygravitational wave.
Key words: gravitational wave, binary black hole, Population III stars

star born as massive star, they possibly become massive
BH. Furthermore, Pop III stars which are the ﬁrst stars
advansed LIGO detected the ﬁrst gravitational signal after the big bang are born as massive star. The typical
GW150914 (Abbott et al. 2016). The source of the gravi- mass of Pop III stars is 10-100 M . Furthermore, thetational wave signal is the binary black hole merger. The wind mass loss is not eﬀective due to no metal. Thus,
Black hole masses of GW150914 are 36 and 29 M (Ab- Pop III star are easier to be massive compact BH. Ourbott et al. 2016). The black holes of GW150914 is more group alredy reseach the gravitational wave from Pop III
massive than conventional black holes (BHs) in X-ray binary BHs and showed that the Pop III binaries typibinaries whose masses are typicaly 10 M (Remillard caly become 30 M binary BHs and some of them merge& McClintock 2006). In order to explain the origin of at the present day. This result predicted the gravitaGW150914, many theories exist such as the Population tional wave events like GW150914 and LIGO paper said
II (Pop II) binary origin, the Population III (Pop III) ʞ recently predicted BBH total masses agree astonishbinary origin, the globular cluster origin and so on. Es- ingly well with GW150914 and can have suﬃciently long
pecially, the Low metal ﬁeld binary origin is widely ac- merger times to occur in the nearby universe (Kinugawa
cepted theory. There are two reason why the low metal et al. 2014) ʟ. In this talk, I talk about why Pop III
ﬁeld binary is widly accepted. First, there are many binaries become 30 M binary BHs and they can merge
massive binaries and the binary fraction of massive stars at the present day.
are high. For example, the binary fraction of O stars
in Milky way young open clusters is 69+− 9% (Sana et al.
2012). Second, if the progenitors of binary black hole are 2. Method
Population I (Pop I) stars, the fraction of massive is few We calculate the Pop III binay BHs and Pop I and II
and they lose a lot of mass due to a stellar wind mass binary BHs for comparison. Pop I and Pop II stars means
loss. Due to strong wind, Population I stars cannot be- solar metal star and metal poor star whose metallicity
come 30 M BH. Furthermore, the binary orbit becomes is less than 10% of solar metallicity, respectively. In thiswide due to the wind mass loss If the progenitor is low talk, we consider four metallicity cases of Z = 0 (Popmetal stars, it possibly becomes massive BH. In the case III), Z = 5 × 10−3 Z , 5 × 10−2 Z (Pop II) and Z = Zof Pop II, they typical mass is same as Pop I, but the (Pop I). There are important diﬀerences between Pop
wind mass loss weaker than that of Pop I. So if Pop II III and Pop I and II. Pop III stars are (1) more massive,
1.

Intriduction

238> 10 M (2) smaller stellar radius compared with thatof Pop I and II (3) no stellar wind mass loss. These
properties play key roles in binary interactions.
In order to estimate the event rate of binry BH mergers and the properties of binary BHs, we use the binary population synthesis code which is the Monte Carlosimulation of binary evolution. First, we choose the binary initial conditions such as the primary mass M1 ,the mass ratio q, the separation a and the eccentricity ewhen the binary is born. These binary initial conditions
are choosed by the Monte Carlo method and the initial
distribution functions such as the initial mass function(IMF), the initial mass ratio function (IMRF), the initial separation function (ISF) and the initial eccentricity
distribution function (IEF). We adopt these distributionfunctions for Pop III stars and Pop I, II stars as Table. 1. Second, we calculate the evolutions of the primary and the secondary stars. If the binary full ﬁlls thecondition of binary interaction, we consider binary interactions such as the Roche lobe overﬂow (RLOF), thecommon envelope (CE) phase, the tidal eﬀect, the supernova eﬀect and the gravitational radiation. We treat
these binary interactions as Kinugawa et al. (2014); Kinugawa et al. (2016) In this paper, we treat the binaryinteraction parameter such as the CE parameter αλ andthe lose fraction β of transfered stellar matter during a
RLOF as αλ = 1, β = 0.
3. Result
We calculate the 106 binaries for each metallicity. Figure
1 is the total mass distribution of BH-BH which merge
within the Hubble time. Black line is the ﬁrst star binary.
The typical total mass of Pop III BBH is 60 M (∼30
M -30 M ). On the other hand, in the other metallicity
cases, the typical total mass is 20 M or so.
There are the reason why Pop III binaries become 30
M -30 M Binary black holes. The Pop III whose massis larger than 50 M evolves as the red giant. The masstransfer of red giant is generally unstable and it becomesthe common envelope phase. Thus, they lose the envelope and they become light. Therefore, the BH mass
become about 30 M . On the other hand, the Pop III
whose mass is less than 50 M evolves as a blue giant.
In such cases, the mass transfer is stable and mass loss
is not so eﬀective. Thus, they become about 30 M BH.
Therefore, the peak of binary BH total mass is about 60
M due to these two evolution path. It does not dependon the IMF and binary parameters. It only depends on
the Pop III stellar evolution. Thus, the peak mass of
Pop III reﬂects the inﬂuence of Pop III stellar evolution.
On the other hand, in the case of Pop I and Pop II, allstars evolve via a red giant, so almost all binaries evolve
via the similar evolution pass. Thus, the shape of Pop II
binary BH mass distribution reﬂect the inﬂuence of IMF

due to similar evolution pass. Pop I has the inﬂuence ofstellar wind due to strong wind.
In this calculation, Pop III binaries tend to become
massive binary BHs. Furthurmore, we can see such binarues which are born at the early universe. Pop III
stars were born and died at z ∼ 10. However, the typical merger time of compact binaries due to gravitationalradiation is too long. Thus, We might see the gravitational wave from ﬁrst stars as binary BH mergers.
In order to calculate merger rate, we need to know
when Pop III stars were born and how many Pop III
stars were born. i.e. we need the star formation rate(SFR) of Pop III stars. We adopt the Pop III SFR by de
Souza, Yoshida & Ioka (2011). The peak value of the
SFR is 10−2.5 M /yr/Mpc3 at z ∼ 9.
Using our population synthesis results and the
SFR, the Pop III binary BH merger rate density
at the present day in our standard model (IMF:
ﬂat 10 M < M <100 M ) is 2.5 ×10−8
events/yr/Mpc3 (SF Rp /(10−2.5 M /yr/Mpc3 ))·([fb /(1+fb )]/0.33). The detection range of KAGRA and advanced LIGO for 30 M -30 M binary BHs is ∼1.5 Gpc.
Thus, the detection rate of Pop III BBH (GW150914
like BBH) in our standard model is ∼ 180 events/yr(SF Rp /(10−2.5 M /yr/Mpc3 ))·([fb /(1 + fb )]/0.33) Thisvalue depends on the SFR and the binary fraction of Pop
III. Since the typical mass of Pop III BBH is 30 M , wecan see quasi-normal mode (QNM) of merged binary BH.
Nakano, Tanaka & Nakamura 2015 show that if S/N of
QNM is larger than 35, we can conﬁrm or refute the General Relativity (GR) more than 5 sigma level. Therefore,
we might not only detect the Pop III binary BH by GW
but also check GR by Pop III BH QNM. Furthermore,
the mass distribution might distinguish Pop III from Pop
I, Pop II. It might become the evidence of the Pop III existence. If the chirp mass distribution can not distinguish
Pop III due to small SFR or so, we can conﬁrm Pop III
BHBH by redshift dependence. DECIGO which is theJapanese space gravitational wave observatory project
has good sensitivity from 0.1 Hz to 10 Hz. B-DECIGO
is test version. DECIGO and B-DECIGO might see suchbinary black holes up to redshift 30. The expected detection rate is about 105 events/yr. The SFR of Pop
III has the peak at z ∼ 10. Thus, it can see Pop III
binary BH when they were born. Furthermore, we cancheck the redshift dependence of high mass binary BH
mergers (See Fig. 2). B-DECIGO can measure the mass
spectrum and the z-dependence of the merger rate to distinguish various models of binary BHs like GW150914,
such as Pop III and Pop II BBH.
References
Abbott, B. P., Abbott, R., Abbott, T. D., et al. 2016,
Physical Review Letters, 116, 061102

239

Table 1: The initial distribution functions.

IMF
IMRF
ISF
IEF

Pop IIIﬂat (10 M <M< 100 or 140 M )ﬂat (10/M < M2 /M1 < 1)logﬂat (amin < a < 106 R )e (0 < e < 1)

Pop I,II
Salpeter (1 M <M< 100 or 140 M )ﬂat (0.1/M < M2 /M1 < 1)logﬂat (amin < a < 106 R )e (0 < e < 1)

Z=0 (Pop III)

Typical total massM䡚60 MÍ
(30 MÍ +30 MÍ)
TK et al. 2014,2016
IMF:Flat(10M<M<140M)

Z=1/200 Zsun

e.g. Pop I, Pop II(Z=0.02,0.001,0.0001)
IMF:Salpeter(1Msun<M<140Msun)
Typical mass 䡚10 MÍ

Z=1/20 Zsun

Z=ZsunTotal mass [Msun]

Fig. 1: The total mass distribution
Abbott, B. P., Abbott, R., Abbott, T. D., et al. 2016,
ApJL, 818, L22
de Souza, R. S., Yoshida, N., & Ioka, K. 2011, A&A,
533, A32
Kinugawa, T., Inayoshi, K., Hotokezaka, K., Nakauchi,D., & Nakamura, T. 2014, MNRAS, 442, 2963
Kinugawa, T., Miyamoto, A., Kanda, N., & Nakamura,T. 2016, MNRAS, 456, 1093
Remillard, R. A., & McClintock, J. E. 2006, ARA & A,
44, 49Sana, H., de Mink, S. E., de Koter, A., et al. 2012,
Science, 337, 444

Log(events/yr)

240

Saturated at z~10

Saturated at z≲5

Pop III BBH

Pop I and II BBH (Belcynski et al. 2016)(2 metallicity evoluon models)

RedshiRedshiFig. 2: The cumlative merger rate of binary BHs.

241

Search for X-ray Counterparts of Gravitational Wave Events
Nobuyuki Kawai,1 Motoko Serino,2 Hitoshi Negoro,3 on behalf of the MAXI Team1

Department of Physics, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8551
2
MAXI team, RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198
3
Department of Physics, Nihon University, 1-8-14 Kanda-Surugadai, Chiyoda-ku, Tokyo 101-8308
E-mail(NK): nkawai@phys.titech.ac.jp
Abstract

The Gas Slit Camera (GSC) of Monitor of All-sky X-ray Image (MAXI) on the ISS scans most of thesky in the 2–20 keV X-ray band every 92 minutes with the orbital revolution of the ISS. Thus it can beused to search for an X-ray counterpart without prompt precise location of the GW event if the emission
lasts for tens of minutes. We examined the MAXI/GSC all-sky X-ray images (2–20 keV) obtained in theorbits preceding and following the gravitational wave event GW150914 and GW151226. In each of the
92-min orbits, MAXI/GSC scanned more than 90% of the localization (90% conﬁdence) regions of the GW
skymap. No signiﬁcant new source was found in these scans with typical upper limit of about 1 × 10−9 erg
s−1 cm−2 in 2–20 keV. Based on this sensitivity, we argue that MAXI may be able to detect afterglows ofshort GRBs within the LIGO/Virgo/KAGRA range for neutron star merger.
Key words: gravitational waves — methods: observational — X-rays: general

1.

Introduction

X-ray/γ-ray emission is expected from a gravitational
wave (GW) source of neutron-star (NS)-NS or NS-blackhole (BH) binaries, while coalescence of BH-BH binary is
not expected to emit electromagnetic (EM) wave. “Short
GRB” has been the prime candidate for EM counterpartof promised GW source. There are some evidence of collimation with opening angle of ∼ 6–30◦ for the emissionsof short GRBs (Fong et al. 2015), and thus the probability of observing jet emissions is only ∼ 4%.
There are several scenarios for soft X-ray production
on the GW sources. Kisaka et al. (2015a) calculated
the light curves with two emission models from thermal
ejecta: one is powered by nuclear decay, and the otheris powered by engine activity such as jet, disk wind, orPoynting ﬂux etc. Kisaka et al. (2015b) suggested thatreﬂected photons from GRB jet can also be a source of
X-ray emission.
The only observation of a prompt emission in 2–25 keVfrom a short GRB is GRB 050709, which was detected
by HETE-2 (Villasenor et al. 2005). Initial short spike of
GRB 050709 was followed by a soft extended emission,
which was observed in 2–25 keV. If short GRBs associatewith GW events and they are followed by soft extended
emissions, MAXI may observe these emissions.

2. MAXI observations of GW sources
MAXI observed the localization regions of two GW
events, GW150914 and GW151226. At the time of theﬁrst event (2015-09-14 09:50:45 UT; Abbott et al. 2016a)
the ISS was about to leave the region with a high particle ﬂux. The MAXI GSC camera was oﬀ at the moment,
but resumed observation 4 minutes later. The second
event occurred at 2015-12-26 03:38:53 UT (Abbott et al.
2016b) when the MAXI GSC was on. Figure 1 shows
MAXI all-sky images observed within an orbit (92 minutes) of the GW events. No signiﬁcant new source wasfound in these scans with typical upper limit of about
1 × 10−9 erg s−1 cm−2 in 2–20 keV. This ﬂux corresponds to the luminosity of ∼ 2 × 1046 erg, if the source
is at a distance of 410 Mpc (Abbott et al. 2016a). The
upper limits for various time scales are shown in table 1.
3. Discussion
We compared the upper limits shown in table 1 withother observations. Connaughton et al. (2016) reported
that Fermi GBM observed a possible gamma-ray transient associated with GW 150914. Assuming the photon
index of −1.4 and using the ﬂux averaged over 1 s, theextrapolated X-ray ﬂux in 2–20 keV is ∼ 2 × 10−8 erg
cm−2 s−1 . An X-ray event with this ﬂux level can bedetected by MAXI GSC even if the emission lasts only
one second.
Another interesting comparison is that with shortFlux (erg/s/cm^2)

242

GRB 050709, which was observed by HETE-2 (the shortpulse and the extended emission; Villasenor et al. 2005)
and Chandra (the afterglow; Fox et al. 2005). Figure 2
shows the MAXI upper limits in three time scales andobserved ﬂuxes of GRB 050709 in the short pulse, theextended emission, and the afterglow phase. The open
symbols are observed ﬂuxes of GRB 050709. The ﬁlled
symbols are those scaled to 100 Mpc, which is the expected distance of the NS-NS merger event for the next
observation run (O2) of LIGO. The ﬁgure shows thatMAXI is capable of detecting X-ray emissions from short
GRBs like GRB 050709, which may associate with GW
events.

10

-6

10

-7

10

-8

10

-9



10

-10

10

-11

10

-12

10

-13

10

-14

10

-15

10
0

GRB 050709 scaled to 100 Mpc

timescale
Flux
Luminosity(s)(erg cm−2 s−1 )(erg s−1 )−10
1 orbit1000< 9.5 × 10
< 1.9 × 1046
4
−10
1 day
8.6 × 10< 2.3 × 10
< 4.6 × 1045
5−1010 days 8.6 × 10
< 0.8 × 10
< 1.6 × 1045
A distance of 410 Mpc is assumed for the luminosity.

10

-5

low

Table 1. Typical upper limits of the MAXI GSC observations in 2–20
keV for various time scales.

-4

g
ter
Af

Fig. 1. MAXI all-sky images observed within one orbit of the GW
events GW150914 (top) and GW151226 (bottom). The thickblack curves show the GW localization regions.

10

-2

10
0

-1

10
0

0

10
0

1

10
0

2

10
0

3

10
0

4

10
0

5

10
0

6

Time since Trigger (sec)Fig. 2. The MAXI upper limits for three time scales () are comparedwith the observed X-ray ﬂuxes of GRB 050709 during the shortpulse (), the soft extended emission, and the afterglow ().
The open symbols are observed ﬂuxes of GRB 050709. The ﬁlled
symbols are the ﬂuxes scaled to 100 Mpc.

4. Summary
Soft X-ray band is unexplored for rapid transients, including possible counterparts of GW events. A large
fraction of GW150914 and GW151226 regions were covered in an orbit, yielding a ﬂux upper limit ∼ 10−9erg s−1 cm−2 . An upper limit for 10 days could be anorder of magnitude smaller. MAXI can constrain the
short GRB scenarios for NS-NS merger at <100 Mpc(O2 range). Instantaneous ﬁeld of view of MAXI is 2%of the sky. It is too small for observing rare GW events.
Observations with instruments of larger ﬁeld of view areneeded.
References
Abbott, B. P. et al. 2016a, Phys. Rev. Lett., 116, 061102
Abbott, B. P. et al. 2016b, Phys. Rev. Lett., 116, 241103
Connaughton, V. et al. 2016, ApJL, 826, L6
Fong, W. et al. 2015, ApJ, 815, 102
Fox, D. B. et al. 2005, Nature, 437, 845
Kisaka, S. et al. 2015a, ApJ, 802, 119
Kisaka, S. et al. 2015b, ApJL, 809, L8
Villasenor, J. S. et al. 2005, Nature, 437, 855

243

X-ray Transient Localization Experiment aboard a micro-satellite formulti-messenger counterpart search
Tatsuya Sawano,1 Daisuke Yonetoku,1 Kazuki Yoshida,1 Yasuaki Kagawa,1
Kentaro Kawai,1 Masao Ina,1 Tatehiro Mihara,2 and Hirokazu Ikeda3
1

Kanazawa Univ., Kakuma-machi, Kanazawa, Ishikawa, 920-1192, Japan
2
RIKEN, 2-1, Hirosawa, Wako, Saitama 351-0198, Japan
3
ISAS/JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa, 252-5210, Japan
E-mail(TS): sawano@ astro.s.
.s.kanazawa-u.ac.jp
Abstract
Short gamma-ray bursts (short GRBs) are likely caused by the coalescence of double neutron star binary
and/or neutron star and black hole binary, and therefore they are expected to be promising electromagnetic
counterparts of gravitational wave (GW) sources. The joint observation of short GRB and GW, and theidentiﬁcation of the host galaxy and distance of such events will provide not only the deﬁnitive observational
evidence for the origins of short GRBs but also the insights to the compact object merger model and the
explosion mechanism. However, GW sky location is uncertain by tens or hundreds of square degrees, which
is not well constrained for the eﬃcient follow-up observation with optical/NIR telescopes. Thus, we aredeveloping X-ray imaging detector that has both ﬁne localization accuracy and wide ﬁeld-of- view, named
Transient Localization Experiment (T-LEX). The energy range, ﬁeld-of- view, localization accuracy of TLEX are designed to be 1 - 15 keV, 1 steradian, and 15 arcminutes, respectively. To achieve such capability,
we adopted two sets of 1-dimensional coded aperture imaging system, which consists of tungsten maskwith random pattern of 1-dimensional strip-like apertures and silicon strip detector with an active area of
50 cm2 for each dimension. T-LEX is plan to be launched onboard a micro-satellite in late FY 2018, which
is being developed at Kanazawa University named as Kanazawa-SAT 3 . The mission life is designed tobe more than 1 year. The objective of the satellite mission is to localize X-ray transients including GW
sources and to send real time alerts of X-ray transient information about the position and the burst trigger
time. We report the imaging performance of a prototype model of T-LEX.
Key words: instrumentation: detectors — micro-satellite - gamma-ray bursts — coded mask imager

1. Introduction
The progenitors and the environment of short duration
gamma-ray bursts (GRBs) (< 2 sec) remain elusive despite of their localizations. To test the theoretical argument that short GRBs are produced in the coalescenceof binary compact objects such as neutron stars or black
holes (Paczynski 1986; Eichler et al. 1989), the observation of short GRBs with gravitational wave (GW) connection gives a decisive evidence.
Recently, the ﬁrst direct GW detection, GW 150914,
was performed by advanced LIGO (Abbott et al. 2016a),and therefore the progress of gravitational astronomy ismuch expected in the future. To maximize the scientiﬁcreturn from the GW observation, it is essential to observe electromagnetic counterparts and identify the hostgalaxies of GW sources. However, the GW detectors
cannot determine the host galaxy of the GW events bythemselves due to their marginal localization ability. For
example, the constraint on the position of GW 150914
was an area of about 600 square degrees (Abbott et al.
2016b). Therefore, the observation of short GRBs witha wide ﬁeld and a ﬁne localization will play a crucial role
to serve a bridge from the GW detection to the electromagnetic follow-up observation.
2. Micro-satellite Kanazawa-SAT 3
Kanazawa-SAT 3 is a micro-satellite with a size of 50
cm cube weighing 50 kg promoted by Kanazawa University, and possesses a wide-ﬁeld X-ray imaging detector(Yonetoku et al. 2014a). The objectives of that satellite mission are to monitor X-ray transients associatedwith GW events, and to provide the position and theburst trigger time of X-ray transients to the groundbased telescope community. Thus, Kanazawa-SAT 3 is244

T-LEX. The energy range of T-LEX is designed to be
1 - 15 keV. The ﬁeld-of-view is more than 1 sr so as tocover the constraint region of GW events eﬃciently. The
localization accuracy is about 15 arcminutes, which iswell constrained for the follow-up observation with optical/NIR telescopes. The eﬀective area of the X-ray
sensors for X-/Y- planes is about 100 cm2 in total.
Table 2. Conﬁguration of T-LEX

Fig. 1. A schematic view of Kanazawa-SAT 3 designed by Meisei
Electric Co., Ltd.

dedicated to contributing to the identiﬁcation of theelectromagnetic counterparts of GW events. Utilizing
the advantage of short development time for a microsatellite, Kanazawa-SAT 3 is planned to be launched inlate FY 2018, in which the global network of the second generation GW interferometers will be in operation.
Table 1 summarizes conﬁgurations of Kanazawa-SAT 3 .
The mission life of the satellite is designed for at least one
year, and three years are set as the operational goal. The
orbit of the satellite is supposed to be a Sun-synchronous
orbit due to multiple launch opportunities.
Table 1. Conﬁguration of Kanazawa-SAT 3

Imaging method
Energy range
Field of view
Localization accuracy
Detector size
Mask aperture fraction

3.2. Design of Coded Aperture System
The coded aperture system of T-LEX consists of a tungsten mask with random pattern aperture and a silicon
strip detector (SSD) with a thickness of 0.5 mm. The
pitches of the SSD and the mask pattern are the same,and 0.3 mm. Generally, the localization pitch of thecoded aperture system for the line-of-sight δθ yields
Mission life
Size
WeightLaunch target
Orbit
Detectordδθ = arctanD> 1 year
3 years (goal)
50 cm cube
< 50 kgin late FY 2018
Sun-synchronouswide-ﬁeld X-ray imaging detector

X-/Y- 1-D coded mask
1 - 15 keV> 1 sr∼ 15 arcmin. (geometrical)
100 cm2 in total (X and Y)
0.5


,

(1)

where d is the pitch of the detector and the mask pattern,
and D is the distance between the mask and detector
planes. Therefore, in order to satisfy the requirementof the localization accuracy, the distance D is set to beabout 70 mm.

Figure 1 shows a schematic view of Kanazawa-SAT .
The satellite has two solar paddles, and a three-axis attitude control system for sun-oriented operation. The
X-ray imaging detector, T-LEX (Transient Localization
Experiment) is installed on the opposite side of the panel
on which the solar panel is mounted, for optical light andthermal protections. Therefore, T-LEX will keep pointing at the anti-solar direction, and can monitor and observe the sky on the night side of the Earth.
3

3. Transient Localization Experiment (T-LEX)
3.1. Detector Overview
T-LEX consists of two (X and Y) sets of 1-dimensional
coded aperture systems. Table 2 shows conﬁgurations of

3.3. Readout Electronics
The charge signals of X-ray events from SSDs are verysmall (∼ 0.04 fC for 1 keV) and therefore they must be
ampliﬁed in order to discriminate with a proper energy
threshold. However, the number of the channels of the
SSD is 512 (X) plus 512 (Y), and one cannot implement
the readout system with discrete circuits. Therefore,
we developed an application speciﬁc integrated circuit(ASIC) for readout of small charge from an X-ray sensor, ALEX (ASIC for Low-Energy X-rays) (Yonetoku etal. 2014b). ALEX has 64 analog inputs and digital input/output serial interfaces. The electrodes of ALEX
inputs and the SSDs are directly connected with bonding wires for least additional stray capacitance. Eachanalog input of the ALEX is processed to a chargesensitive ampliﬁer, two shaping ampliﬁers with slow and245

same. We made a case of the prototype of T-LEX madeof aluminum, where (d/D)−1 is designed to be 2.24×102.

Fig. 2. A picture of the sensor boards stacking on the readout boardsof the T-LEX prototype model.

fast time constants of several microseconds, and a comparator from the fast shaping ampliﬁer output, and ananalog-digital converter from the slow shaping ampliﬁer.
A ﬁeld programmable gate array (FPGA) is used tocontrol the ALEX and process the event data throughthe digital interfaces of the ALEX. In this way, eight
ASICs are used to readout the X- (or Y-) sensor andthose ASICs are controlled with one FPGA. To control
two (X- and Y- ) FPGAs, another FPGA is used for
merging the X- and Y- event data and communication
with the CPU on the satellite bus system. In the backend readout FPGA, the logic of one CPU core is implemented so as to manage complicated procedures suchas a burst trigger system and the serial communication
with the bus system.
3.4. Prototype Model
We developed a prototype model of T-LEX. A pictureof the sensor boards stacking on the readout boards of
the T-LEX prototype model is shown in Figure 2. The
SSD installed on the sensor board in the prototype modelcomprises 512 strips for each (X- and Y-) dimension, andeach strip size is 0.3 × 16 mm2 . Therefore, the detector
area of the SSD for each dimension is ∼ 25 cm2 , whichis the half of the full conﬁguration at the ﬂight model.
On the other hand, the number of the readout channels is the same as that of the ﬂight model, and thus
this prototype model can be characterized as the engineering model of the T-LEX electronics. There are eight
ASICs on each (X- and Y-) sensor board, and the FPGA
is mounted on the readout board on the back of the sensor board. The FPGA with the implemented CPU isalso mounted on the X- readout board. To reduce thechanges in the geometrical and functional designs for theﬂight model, the almost of the all electronics are pincompatible devices with the space-grade ones.
As a trial piece of the random pattern mask, we madea small tungsten mask with a thickness of 35 microns.
The aperture size of that tungsten mask is 76.8 × 76.8mm2 for one dimension, which is a quarter of the ﬂight
model. We note that we cannot perform the full-coded
imaging test using the all channels of the SSDs because
the widths of the sensor and the mask aperture are the

4. Imaging Performance of T-LEX Prototype Model
4.1. Overview
To conﬁrm the validity of the basic design of T-LEX,
we performed an imaging demonstration test with partial channels of the SSDs. The goal of that test is asfollows. First, we measure the reconstructed directionsof X-ray beams with various incident angles. Second, we
compare those experimental data with the ones obtained
geometrically. Last, whether the discrepancy of them is
acceptable or not is discussed.
4.2. Reconstruction method
To obtain the reconstructed image of the coded mask
system, we used a cross correlation method expressed
with
rj =

di · mi+j ,(2)iwhere the j is the shift index in the cross correlation
calculation, rj is the intensity at the shift index of jin the reconstructed image, di is the number of the Xray events at the detector position of i, and mi+j is themask weighing design element. In this paper, we usedthe condition on the mask weighing where mi is 0 if themask is opaque at i and 1 if the mask is transparent.
4.3. Setup of Experiment
As the parallel light of X-ray photons, we used a
5.5-meter beam line at the astrophysics laboratory ofKanazawa University. X-ray photons from an X-ray generator with a tube voltage of 40 kV irradiate the molybdenum target with a size of 5 × 5 cm2 so as to produce ﬂuorescent X-rays with an energy of 17.5 keV (Kαline). The T-LEX prototype model is placed in a thermostatic bath, which is at a distance of 5.5 meters fromthe molybdenum target with an ambient temperature of−10 degrees Celsius. A rotation stage is installed on
the side of the case of the T-LEX prototype model inorder to change the incident angle of the X-ray beam.
The divergenceof the X-ray beam is calculated to be√(5 cm/ 2)/(5.5 m) ≈ 0.0064, or 22 arcminutes, and it isslightly larger than the localization accuracy of T-LEX.
This eﬀect causes the peak dilation in the reconstructed
image.
As mentioned in Section 3.4., we cannot use the wholechannels under the full-coded condition. Thus, we usedthe data from two ASICs, or 128 channels that are close
to the center. Using those 128 channels of X-ray intensity
data and 256 strips of the mask patterns, we obtained thereconstructed image according to Eq. 2. We obtained
the peak position of the reconstructed image by ﬁttingthe data around the peak with a normal distributionfunction.
4.4. Result and Discussion
Figure 3 and 4 show examples of di and rj , respectively,
obtained with the experiment. Figure 5 shows the peak
position in terms of the shift index versus the X-ray incident angle (or the rotation angle), where the origin isarbitrary. Ideally, the curve can be modeled with a tangent function because the peak position m yields
md,D

4.0×105

3.5×105

3.0×105

model
measured

2.5×105
2.0×105

1.5×105

1.0×105

5.0×10

4

0

50

100
150
Shift Index

200

250

(3)

280

model fit

230
Shift Index

where θ is the rotation angle from the line-of-sight. We
ﬁtted this curve with a tangent function and obtained(d/D)−1 to be 2.27 × 102 using Eq. 3, which is 1.5%
larger than that designed. That means if we use thedesigned value for (d/D)−1 , the reconstructed angle forthe events with an incident angle of 30 degrees has anerror of about 30 arcminutes, which is not acceptable dueto the requirement. Therefore the calibration of (d/D)−1
by the imaging test will work eﬀective to compensate
that systematic error. It is suggested that the imaging
test will be mandatory also for the ﬂight model.

Fig. 4. An example of the one-dimensional reconstructed image
within the half-coded region (red solid line). The model image
for parallel light is also drawn (gray dashed line).

180

Residualtan θ =

Correlation Intensity (arb. unit)

246

80
0.3
0
-0.3

130

-20

-10

0

10

20

Incident Angle (degrees)Counts (arb. unit)

6000

Fig. 5. The peak positions in terms of the shift index in reconstructed
image versus the incident angles obtained by the experiment (redcross). The coordinate of the origin is arbitrary. The best estimateof the tangent model function is drawn in the blue solid line. The
residuals are also shown.

4000

2000

0
80

100

120
140
Strip Number

160

180

Fig. 3. An example of the X-ray intensity distribution (yellow-ﬁlled
histogram) representing the mask pattern obtained by the X-ray
beam irradiation. Behind the histogram, the best matching position of the mask is shown, where the transparent elements of themask are represented with the gray color region.

5. Summary
The X-ray transient localization with a wide ﬁeld-of-view
will strongly promote the electromagnetic follow-up observation for the identiﬁcation of the host galaxy of GW
events, contributing to the progress of GW astronomy.
The micro-satellite Kanazawa-SAT 3 possesses the wideﬁeld X-ray imaging detector, T-LEX, and the objectivesof that satellite mission are to monitor the X-ray transients associated with GW events and to alert to theground-based telescope community sending the information of the position and the burst trigger time of X-ray
transients.
We developed the prototype model of T-LEX and successfully performed the imaging test. We found that theimaging performance test eﬀectively works as the calibration for reconstructing the incident angles. That calibration will be also performed carefully to the ﬂight modelof T-LEX to satisfy the requirement for the localization
accuracy.
References
Abbott, B.P. et al., 2016, Phys. Rev. Lett., 116, 061102
Abbott, B.P. et al., 2016, Phys. Rev. Lett., 116, 241102
Eichler, D. et al., 1989, Nature, 340, 126
Paczynski, B. 1986, ApJ, 308, 43
Yonetoku, D. et al., 2014, UNISEC Sp. Takumi J., 5, 19
Yonetoku, D. et al., 2014, SPIE, 91442S

247

Explore the Transient X-ray Sky with Einstein Probe
Weimin Yuan
Key Laboratory of Space Astronomy and Technology, National Astronomical Observatories,
Chinese Academy of Sciences, Beijing 100012 China
E-mail(WY): wmy @ nao.cas.cn
Abstract
The Einstein Probe is a small mission dedicated to time-domain astronomy to monitor the sky in thesoft X-ray band. It will carry out systematic survey and characterisation of high-energy transients at unprecedented sensitivity, spatial resolution, Grasp and monitoring cadence. Its wide-ﬁeld imaging capability
is achieved by using established technology of micro-pore lobster-eye X-ray focusing optics. Complementary to this feature is deep X-ray follow-up capability enabled by a narrow-ﬁeld X-ray telescope with alarger eﬀective area. It is also capable of real time triggering and downlink of transient alerts on the ﬂy, inorder to activate multi-wavelength follow-up observations by other astronomical facilities worldwide. Itsscientiﬁc goals are concerned with discovering new or rare types of transients, particularly tidal disruption
events, supernova shock breakouts, high-redshift gamma-ray bursts, and possibe electromagnetic sources
associated with gravitational wave events.
Key words: X-ray: instruments — X-ray: time-domain astronomy — mission: Einstein Probe

1.

Introduction

The X-ray sky is rich in transients and variables of various types. With diverse timescales from sub-seconds toyears, a large variety of such dramatic objects have been
discovered and extensively studied ever since the earlydays of X-ray astronomy, thanks to successive all-sky
monitors in the X-ray (as well as γ-ray) waveband. In
recent years the successful operations of Swift (Gehrels et
al. 2004) and MAXI (Matsuoka et al. 2009) have greatly
expanded our horizon in monitoring the X-ray sky andadvanced our knowledge about the dynamic Universe(Gehrels & Cannizzo 2015; Mihara et al. 2017). New
transients, particularly of previously unknown and scientiﬁcally important types, continue to be discovered.
These include new Galactic black hole X-ray transients(candidates), relativistic outbursts on long timescales interpreted as arising from jets launched by massive black
holes following tidal disruption events, gamma-ray bursts(GRBs) at high redshifts up to z ∼ 9, supernova shock
breakouts, etc. To characterise and understand thesenew phenomena, a large sample of events and detailed
observations are needed.
With the advent of major wide-ﬁeld sky-monitoring facilities across the entire electromagnetic spectrum, such
as LSST and LOFAR/SKA, the turn of the decade will
see a golden age of time-domain astronomy with ﬂourishing discoveries, as generally anticipated. Even more
excitingly, gravitational-wave (GW) astronomy will mature with the 2nd generation of GW detectors being able

to detect a potentially large number of events. They callfor synergy observations across the electromagnetic-wave
spectrum to detect their potential electromagnetic counterparts. This is essential for identifying the associatedobjects/galaxies and to try to comprehend the nature of
GW events.
In X-ray, the above driving science invokes the next
generation of instruments with higher sensitivity andimproved angular resolution (a few arc-minutes or less)than those currently available. These requirements are
now readily enabled by novel X-ray focusing optics—thelobster-eye micro-pore optics, in which X-ray focusing
imaging results in enormously enhanced gain in signal
to noise, and thus high sensitivity, while a wide Fieldof-View (FoV) can be maintained. The Einstein Probe
mission (Yuan et al. 2015) is based on such a technology,
with an aim to monitor a large fraction of the sky athigh cadences with unprecedented sensitivity in the soft
X-ray band.
2. Scientiﬁc objectives
The Einstein Probe (EP) is a short-listed mission of priority in the Space Science Programme of the Chinese
Academy of Sciences (CAS) for the nation’s 13th ﬁveyear plan. Its aims are to discover and characterise
high-energy transients and monitor variable objects inthe soft X-ray band with sensitivity one order of magnitude higher than those currently in orbit.
The key scientiﬁc objectives of the Einstein Probe mis-

248sion are:(1) Carry out systematic all-sky surveys of high-energy
transients in the soft X-ray band at high cadence overa wide range of timescales from seconds to years with
unprecedented sensitivity(2) Discover quiescent black holes at almost all astrophysical mass scales and study how matter falls onto
them by detecting transient X-ray ﬂares, particularlystars being tidally-disrupted by otherwise dormant massive black holes at galactic centres.
(3) Discover the X-ray photonic counterparts ofgravitational-wave transients found with the next generation of gravitational-wave detectors and precisely locate
them.
The mission will address some of the key questions inastrophysics and cosmology, such as the prevalence of
black holes in the Universe and how they formed andevolved, as well as their interplay with host galaxies;
how black holes accrete mass and launch jets; the astrophysical origins and underlying processes of gravitational
wave events; details of the physics which operates in extreme conditions of strong gravity; when and where did
the ﬁrst stars form in the early Universe and how they
re-ionise the Universe; the progenitors and processes ofsupernovae (Yuan et al. 2016a).
2.1. Systematic census of transients in the soft X-ray sky
EP will carry out systematic surveys to discover andcharacterise high-energy transients of various types overa wide range of timescales and at high cadence in thesoft X-ray band. It will also perform immediate followup observations of newly discovered transients with its
narrow-ﬁeld X-ray telescope onboard, and will issue fastalerts to trigger follow-up observations by global multiwavelength facilities. Most of the previous and current
high-energy astrophysical monitors are operating in themedium and hard X-ray (even gamma-ray) bands, andthe soft X-ray sky (from a fraction of 1 keV to several
keV) has hardly been monitored extensively with both
suﬃcient sensitivity and cadence. It is thus expectedthat EP will open up a new window in time-domain Xray astronomy. This implies that the bulk of the targets
and processes being monitored will probably change fromthe non-thermal, highly relativistic beaming regime (e.g.
GRBs, blazars) to thermal, low-/non-beaming regime(e.g. accretion, shock), and thus new types of transientsare expected to emerge. Above all, its increased sensitivity will enable unprecedented census of faint or distanthigh-energy transients beyond the horizons of the previous and contemporary X-ray monitors. As such, a large
variety of X-ray transients are expected to be detected,
even possibly previously unknown types.
Of particular interest, EP is expected to discover alarge sample of black hole tidal disruption events, outnumbering the sparse and heterogeneous cases currently

known, and may revolutionise the research of the ﬁeld
(see next subsection). Other important types of transients include GRBs at higher redshifts than the currently known sample (7 GRBs with z=6.2–9.4). These
valued events, produced by violent death of stars in theearly Universe, would carry unique information on earlystar-formation and metallicity evolution, and hopefullyon the ﬁrst generation of stars and the re-ionization in
the dark early Universe, which are otherwise almost inaccessible from ground-based facilities (Yuan et al. 2016b).
Shock breakout emission from supernovae is the prompt
X-ray emission produced as the outward-propagating
shocks generated by core-collapse breaks out of the stellar surface. X-ray observations can yield important clues
to the properties of the progenitor stars. There are only
a few candidate events known so far due to their elusiveness given the short burst durations (∼ 103 s) andmoderate brightness. EP is also expected to detect moreof such events. Other transient sources to detect andcharacterise in large numbers include X-ray ﬂashes, lowluminosity GRBs, X-ray rich GRBs, GRB precursors,
magnetars, stellar corona ﬂares, classical novae, supergiant fast X-ray transients, and outbursts of active galactic nuclei and blazars, etc.
2.2. Tidal disruption events and quiescent black holes
Stars are tidally disrupted and accreted when their selfgravity cannot balance the strong tidal force as they approach massive black holes (MBHs) closely enough, producing a ﬂare of electromagnetic radiation, preferably
peaked in the UV and soft X-ray regime as predictedin theoretical models (Rees 1988). First discovered in
the ROSAT survey (Komossa & Bade 1999), there are
only a few dozen TDE candidates found so far, mostly inthe declining phase via searches from multi-wavelengtharchives and surveys (Komossa 2015). As perhaps themost unique signature of the existence of MBHs in otherwise quiescent galactic nuclei, a systematic survey ofTDEs provides a census of quiescent MBHs in the Universe, i.e. to constrain their occupation fractions in various types and masses of galaxies, which is essential forunderstanding the formation and evolution of massive
black holes and of galaxies. Recent discoveries of a few
jetted events by Swift and MAXI indicated that relativistic jets can be launched in at least some of the TDEs
(Burrows et al. 2011). A distinctive X-ray lightcurvewith large drops and recurrences was also seen in one
case, which was interpreted as perhaps the most accessible signature of quiescent binary MBHs, via tidally disrupting a star (Liu et al. 2014).
With its soft X-ray bandpass and high sensitivity, EP
is an ideal instrument to discover and study TDEs ina systematic way, and to catch them at the peaks of
their X-ray ﬂares out to at least a few hundred Mpc.
The estimated detection rate of TDEs may range from

249

several tens to hundreds per year, depending on the actual event rate per galaxy and the occupation fractionof MBHs. The resulting homogeneous sample of TDEs
obtained by EP in several years’ operation will provide
useful constraints on the galactic occupation fraction ofMBHs, by combining the incidence of TDEs for various
type of galaxies, which can be determined from simulations. Some of the TDEs may be caught at the risingphase of the ﬂares, making it possible to observe at multiple wavelengths over the entire cycle of TDEs from theonset of disruption till complete fading away.
Moreover, EP is expected to discover new stellar- andintermediate-mass black holes lurking in our Milky Way
and in nearby galaxies, by detecting their X-ray outbursts due to some kind of instability of gas accretion.
In this sense, EP is a black hole ﬁnder at almost all astrophysical mass scales.
2.3.

Detecting electromagnetic sources of gravitationalwave events
Gravitational wave (GW) astronomy has come of dawn
with the direct detections of the ﬁrst GW event ever,
GW150914 (Abbott et al. 2016) and subsequent a few
others. Up on the completion of upgrade of Advanced
LIGO, Advanced Virgo and other GW observatories, it
is highly anticipated that, at the turn of the decade, GW
events will be detected in large numbers. The detectionof their electromagnetic (EM) counterparts is of greatscientiﬁc signiﬁcance in the sense that astrophysical observations in the EM domain will provide independent
evidence for further establishing the model interpretation of the events solely based on the general theory
of Relativity. Moreover, as the loci of GW sources canonly be determined within regions from several hundreds
to several tens of square degrees even for the upgradedglobal GW detector network, the detection of EM counterparts is essential for precisely locating the GW eventsand thus the identiﬁcation of associated astrophysicalobjects. This will enable precisely measuring redshifts(thus the source distance and EM energy budget) andhelping understanding the nature, origin and evolutionof the progenitor systems. As a potentially even more
far-reaching cosmological application, joint detections of
GW and EM sources (redshifts) can be used to probe
the geometry and expansion of the Universe, and to determine the cosmological parameters (Schutz 1986).
Although no EM counterparts have been reported forthe a few GW detections made so far, strong EM radiation associated with mergers of binary neutron stars
or neutron star–black hole binary has been predicted(Zhang 2013). With the combined high sensitivity, large
ﬁeld of view, manoeuvrability and rapid alert downlink,
EP has great potential in detecting the EM counterparts of GW events by synergy with the next generationof GW detectors. This may be achieved either by di-

Fig. 1. Illustration of light path of a point-like source in a lobster eye
optics (Credit: adopted from http://www.x-ray-optics.de/).

rectly detecting X-ray transient sources associated with
GW events that happen to be within the large EP FoV,
and/or by rapidly observing the sky region to cover thelarge loci area provided by the GW detectors as ToO observations. For the latter, fast uplink of GW alert data
is essential.
3. Lobster-eye micro-pore optics
In general, X-ray focusing is realised by multiple grazing incidence refections on surfaces almost parallel tothe direction of incident X-rays. The conventionally andcommonly adopted conﬁguration is the Wolter-I optics,
which reﬂects X-rays by a tubular, rotationally symmetric parabolic surface followed by a hyperbolic surface. The FoV of such optics is inherently small, typically less than one degree. Alternatively, the Lobster-eye
Micro-Pore Optics (MPO) (Angel 1979), which mimics
the imaging principle of the eyes of lobsters (Figure 1),
can reach an unrestricted FoV. The lobster-eye optics ismade of a thin spherical micro-channel plate with millions of square micro-pores, the axes of which all point
radially to a common centre of curvature. Incoming lightat a grazing incidence angle is reﬂected oﬀ the walls of
the many tiny pores, and is brought onto a focal spherewith a radius of half the curvature of the optics. For any
given direction of incoming light, a same large numberof pores, whose reﬂection surfaces are conﬁgured orthogonal to each other, take part in the reﬂection. Such a
conﬁguration has no preferential optical axis, and canthus deliver an un-vignetted FoV as large as, in principle, the entire solid angle of 4π.
Such a focusing optics produces true imaging witha cruciform point spread function, which has a central
peaking spot produced by two reﬂections from adjacentwalls of the pores and cross-arms produced by single reﬂections. It can yield a large focusing gain of ∼2000,
and a moderate resolution of FWHM the order of arcminutes delivered by the currently commercially available MPO pieces. In practice, the MPO Lobster-eye optics can achieve a FoV of thousands of square degrees

250

Fig. 2. A demonstration prototype of a lobster-eye MPO telescopeassembly developed at X-ray Imaging Lab, NAOC (left-hand), anda true X-ray image formed on its focal plane for a point-like source(right-hand) showing the characteristic cruciform PSF of the lobster-eye optics (credit: C. Zhang @ NAOC, CAS).

with very light weight, which is unique for wide-ﬁeld Xray imaging. Figure 2 shows a demonstration prototype
of a portion of the optics assembly of one WXT module(left panel), as well as an X-ray image obtained at the
focal plane for a point-like X-ray source, i.e. PSF (rightpanel).
4. Mission concept
4.1. Instruments
EP will carry two scientiﬁc instruments—an X-ray monitoring instrument Wide-ﬁeld X-ray Telescope (WXT)
with a large instantaneous FoV, and a narrow-ﬁeld(∼ 30 ) Follow-up X-ray Telescope (FXT). To achieveboth wide FoV and X-ray focusing, the novel micropore optics in the lobster-eye conﬁguration (see above) is
adopted for WXT. WXT consists of 12 identical modules
with a 375 mm focal length, each of which is composed of
6 × 6 mosaicking MPO plates, subtending ∼ 300 square
degrees. Figure 3 (upper panel) shows the layout of one
WXT module. An optical baﬄe is attached at the front
end of the MPO assembly to shield optical stray lightfrom the Sun, the Moon and the Earth. The 12 modules make a total vignetting-free FoV of WXT of ∼ 3600
square degrees (∼1.1 steradian), as illustrated in Figure 3(lower panel).
WXT has a large-format focal plane of approximately
420 mm×420 mm. The baseline choice of the focal plane
detectors is CMOS imaging sensors, developed at the
Gpixel Inc. in China. The current design is to mosaic the
focal plane detector of one WXT module by 2x2 CMOS
sensors. This new type of detectors, which are currently
being under development, have 4 k×4 k pixels (of pixel
size 15 micro) and 6 cm×6 cm in size. The CMOS detectors have some advantages over CCDs for their fast readout speed, as fast as several tens of frames per seconds.
They can thus be operated at moderately low temper-

Fig. 3. Upper panel: Design of one module of the wide-ﬁeld X-ray
telescope. Lower panel: Illustration of the ﬁeld of views of the
WXT modules and the follow-up X-ray telescope (credit: C.
Zhang @ NAOC, CAS).

atures. One WXT module weights 17 kg including the
MPO mirror assembly, detector and electronics unit, optical baﬄe, structure and thermal control.
The nominal detection bandpass of WXT is 0.5–
4.0 keV. Figure 4 shows the eﬀective area curves of WXT
derived via realistic ray-tracing simulations taking into
account the imperfectness of the MPO arrays (Zhao et
al. 2014). The focal plane detector is a CMOS sensor
coated with 200 nm-thick Aluminum. The peak eﬀective
area is ∼3 cm2 at around ∼1 keV for almost any direction
within the FoV except at the edge of the MPO plates.
Thought the eﬀective area for a given direction is nothigh, it has nearly the same value across the entire FoV,
yielding a very large Grasp (eﬀective area times FoV)of the order of ∼ 104 cm2 sq.deg., which is the largest
among all focusing instruments in X-rays. Based on
simulations of the X-ray background on the detector, including particle background and incident diﬀuse X-ray
emission (both the cosmic X-ray background and Galactic diﬀuse emission), we estimate a detecting sensitivity
for a point-like source1 for WXT to be approximately*1 Assuming a power law spectrum with a photon index of -2 anda Galactic absorption column 3 × 1020 cm−2 .

251

Table 1. Speciﬁcations of WXT and FXTParameters
Number of modules
Field-of-view
Spatial resolu. fwhm
Bandpass (keV)
E resolution (eV) @1 keV
Eﬀective area @1 keV (cm2 )Fig. 4. Simulated eﬀective area curves of WXT for the central
focal spot and plus the cruciform arms. The MPO arrays arecoated with Iridium, and have surface roughness of 0.55 nm
and the tilts of pores following a Gaussian distribution with
σ=0.85 arcmin. The focal plane detector is a CMOS sensor coatedwith 200 nm-thick Aluminum. For details of the simulations see
Zhao et al. (2014) (credit: D. Zhao @ NAOC, CAS)

∼ 1 × 10−11 ergs s−1 cm−2 (in the 0.5-4 keV band) or∼0.3 mCrab at 1000 second exposure2 , much more sensitive than previous and current wide-ﬁeld X-ray monitors,
including MAXI and Swift/BAT. The angular resolutionis ∼5 arcmin (FWHM) for the central focal spot, limited by the manufacturing precision of the MPO plates
for the current technology, and possibly by misalignmentand distortion of the plates introduced in the process ofassembling.
The EP FXT is composed of a focusing mirror assembly, which is of the conventional Wolter-I type with a
focal length of 1.6–2 m (TBD) and a CCD detector at
the focal plane, cooled to a temperature in the range of-70◦ – -100◦ C . Its eﬀective area, ≥ 120 cm2 at around
1 keV, is much larger than that of WXT, covering an energy passband of 0.3–8 keV. The total weight of the FXT
module is about 100 kg.
The speciﬁcations of WXT and FXT are summarisedin Table 1. The combination of an array of wide-ﬁeld
lobster-eye modules to act as a soft X-ray transient monitor and a narrow-ﬁeld telescope for deep follow-up observations gives unprecedented Grasp and sensitivity thatwill revolutionise X-ray time-domain astronomy.
*2 It should be noted, however, that these values are based on
the simulations and may be considered to be nominal; thetrue sensitivity may be somewhat lower considering possible
degradation of the performance of the actual MPO plates used,
as well as the actual environment of charged particles in orbit.

WXT

FXT

12
3600 sq.deg.
5’
0.5–4.0
170
∼3

1
∼ 30
<2’
0.3–8.0170≥120

4.2. Mission ProﬁleOne possible conﬁguration of the EP payload is shown in
Figure 5. The payload has a weight of 340 kg and powerof 430 W, and whole satellite has a wight of ∼900 kg andpower of ∼800 W in total. The satellite platform is tobe provided by MicroSat, CAS.
The satellite will be in a circular orbit at an altitudeof 600-650 km and a period of ∼97 minutes, and an inclination angle of ∼30◦ . The survey strategy of EP will
be composed of a series of pointings to mosaic the night
sky in the directions avoiding the Sun, with an avoidance angle to the edge of the WXT FoV always greaterthan 90 degrees, as illustrated in Figure 6. During its 97minute orbit of the satellite, three ﬁelds will be observedwith WXT on the night-side of the sky, each with a 15–
20 min exposure. Over three successive orbits almost theentire night sky will be covered, with sampling cadences
for a given sky region ranging from 5 to 25 times perday. Monitoring the night sky also makes it possible toperform follow-up observations of detected transients byground-based optical/IR telescopes. The set of pointing
directions is shifted by about 1 degree per day to compensate the daily movement of the Sun on the sky. In thisway, the entire sky will be covered within half-a-year’s
operation of the mission.
There are three basic observing modes for EP: themonitoring survey mode using WXT, the following observation mode using mainly FXT (while WXT continues to be in data acquisition mode), as well as the ToO
observation mode of both WXT and FXT upon requests
from the ground segment via the command data uplinkroute.
During observations of WXT, onboard computers will
search for transients from X-ray photon events collected
in real time over a range of timescales. Once a transient
source is detected, classiﬁed, and triggered by the dataprocessing and alerting system onboard, the satellite will
slew (with a slewing speed of 15 degrees per minute) toa new position to enable pointed follow-up observations
with FXT by targeting the new source within its FoV.
Meanwhile, WXT continues to monitor the new sky re-252

Fig. 6. Illustration of the ﬁeld-of-view and pointed observations in
one orbit (credit: MicroSAT, CAS).

Fig. 5. A possible conﬁguration of the EP payload, with twelve WXT
modules and one FXT module (at the centre) (credit: MicroSAT,
CAS).

gion centering the position of the transient. It is essential to transmit alerts quickly to the ground segment,
preferably within one minute or so. The baseline design
is to make use of the short text message capability ofBeidou—the Chinese satellite navigation system, whichcurrently covers the Asia-Paciﬁc region only and will be
extended to complete the global coverage by 2020. One
interesting feature of the Beidou system is its quick uplink capability, which will enable fast ToO observations,as required by follow-up observations of urgent transient
sources such as GW events. An alternative possibility
is also under investigation to make use of the French
VHF ground station system, which is being built for thejoint Chinese-French GRB mission SVOM. The normal
telemetry data will be transmitted to the ground segment via the X-band using the ground tracking stationsof the CAS, mainly at Sanya, Hainan in China.
The mission, aimed for launch around 2022, is designed to have a nominal lifetime of three years witha goal of ﬁve years. The consortium of EP payload is led
by Key Laboratory of Space Astronomy and Technology,
National Astronomical Observatories, CAS, in collaboration with a range of institutions in China and abroad,including Institute of High Energy Physics, Shanghai
institute of Technology Physics, University of Leicester(UK), etc. The Space Science Programme of the CAS
is managed by the CAS’s National Space Science Centre(NSSC).
The author would like to thank the conference organisers
for invitation to this interesting meeting. Gratefulness

is given to the members of the Einstein Probe team,
who have had helped providing the materials and content
of this paper. The work is supported by the Strategic
Priority Research Program of the Chinese Academy of
Sciences, Grant No. XDB23040100.
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255

Development of the Hard X-ray Monitor onboard WF-MAXI
Satoshi Sugita,1 Shohei Harita,1 Yoichi Yatsu,1 Naoki, Isobe,1 Nobuyuki Kawai,1 Makoto Arimoto,2Hirokazu Ikeda,3 Hiroshi Tomida,3 Shiro Ueno,3 Tatehiro Mihara,4 Motoko Serino,4Hiroshi Tsunemi,5 Atsumasa Yoshida,6 Takanori Sakamoto,6 Takayoshi Kohmura,7Hitoshi Negoro,8 Yoshihiro Ueda,9 Yohko Tsuboi,10 and Ken Ebisawa, 31Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8550
2
Waseda University 3-4-1 Okubo, Shinjuku-ku, Tokyo, 169-8555
3
ISAS, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa, 252-5210
4
RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198
5
Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043
6
Aoyama Gakuin University, 5-10-1 Fuchinobe, Chuo-ku, Sagamihara, Kanagawa 252-5258
7Tokyo University of Science, 2641 Yamazaki, Noda, Chiba 278-8510
8
Nihon University, 1-8-14 Surugadai, Chiyoda-ku, Tokyo 101-8308
9
Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502
10
Chuo University, Kasuga 1-13-27, Bunkyo-ku, Tokyo 112-8551
E-mail(SS): sugita@hp.phys.titech.ac.jp
Abstract
Wide-Field MAXI (WF-MAXI) is a mission to detect and localize X-ray transients with short-term
variability including EM counterparts of gravitational-wave events such as gamma-ray bursts and supernovae etc. WF-MAXI consists of two main instruments, Soft X-ray Large Solid Angle Camera (SLC) for
localization of the events with few arc-minutes error in soft X-ray band and Hard X-ray Monitor (HXM)
to measure the spectrum and the lightcurves in hard X-ray band. The ﬁeld of view of WF-MAXI is 20%of all sky at any given time and the energy band is covered from 0.7 keV to 1 MeV. We have developed
the HXM components which consists of 24 channel arrays of Ce:GAGG scintillators coupled with APDs
covering the hard X-ray band with eﬀective area of above 120 cm2 . To read out signals from the censor
arrays, we designed a new LSI dedicated with the readout of 32 APDs’ signals using 0.35 μm CMOS
technology. We will report the integration and the end-to-end test of the components of HXM BBM.
Key words: Instrumentation: X-ray detectors – APD – VLSI – WF-MAXI

1.

Hard X-ray Monitor (HXM)

Hard X-ray Monitor (HXM) is a transient monitor
for electromagnetic counterparts of Gravitational-Wave(GW) events. It was designed for WF-MAXI which
was a mission for detection of X-ray emission fromgravitational-wave sources with short-term variability
and localization of the sources to be enable for follow-up
observation in early time. HXM is designed for wideband spectroscopy with high-time resolution for identiﬁcation of the EM counterparts such as short GRB.
HXM consists 1-D multi arrays of Ce:GAGG scintillator
coupled with avalanche photo-diodes (APDs) with wide
FoV: ∼25% of the whole sky (Fig.1). From a requirementof the mission, the detector is small size of ∼15×15×10
cm3 and total power consumption of <10 W. The energy band is 20 keV – 1 MeV with an eﬀective area of
120 cm2 at 30 keV. HXM also has localization capability

to utilizes a shaded pattern created by passive shields
with accuracy of a few degrees.

Fig. 1. Schematic view of HXM

2. Development and Performance of VLSIWe have developed new VLSI dedicated for ampliﬁcation
and readout of signals from APDs, which has Wilkinsontype analog-digital converters (e.g., mixed signal LSI).

256

We utilized a well-studied 0.35 μs CMOS technology;
Open IP project initiated by Prof. Ikeda and accumulated knowledge based on developments of CdTe sensors(Sato et al. 2011).
To clear the requirement of the lower limit of energy
band of 20 keV, ampliﬁers of the VLSI must be low noise.
However, as a detector capacitor of APDs is very large(∼100 pF), the ampliﬁer noise is very crucial for detection of the low energy photons. In this case, the noise iscomposed of Johnson and ﬂicker noises. To reduce thesenoises, we then designed an input transistor of charge
sensitive ampliﬁer (CSA) to make a drain current larger
and the parallel number of the input transistor larger.
Furthermore by transient and AC analyses using SPICE
simulator, we determined ﬁnely tuned circuit parameterswhich fulﬁlled low noise of ∼2100e- (RMS) at 100 pF.
Then we fabricated the VLSI with function of 32-ch analog ampliﬁers and AD converters optimized for APDs.
From the measurement of the fabricated VLSI, the waveforms (Fast and Slow) were almost the same as the simulated one (Fig.2) and noise performance achieved 1930e- +1.9e-/pF (RMS).

and the noise peak was lower than the lower limit ofenergy band of 20 keV. The energy resolution was 54.3%
at 32 keV and 7.2% at 662 keV.

Fig. 3. Size of a unit of crystal and APDs for scintillator-arrays.

Fig. 4. The spectrum of 137Cs measured by Ce:GAGG+APD and
new VLSI. The peak near 0 keV is noise.

Fig. 2. The experiment and simulated waveforms of Fast and Slowsignal by new VLSI for HXM

3. APDs and GAGG crystals
A unit of scintillator arrays consists of scintillator crystal
with size of 50×10×10 mm3 coupled with APD with sizeof 10×10 mm2 (Fig.3). 24 units are arranged in 2 rows
(12×2) on aparture. We adopted Ce:GAGG (Ce-doped
Gd3 Al2 Ga3 O12 ) crystal which had advantages; no hygroscopic, large light yield of 46,000 photons/MeV and fast
decay time <100 ns. We also adopted reverse-type APDsof 5×5 mm2 (S8664-55 provided by Hamamatsu Photonics) which are ﬂight-proven technologies developed by
TSUBAME (Yatsu et al. 2014) and ASTRO-H (Kataoka
et al. 2012).
We measured the spectrum using the scintillator unit
and the new VLSI. Figure 4 shows the spectrum of 137Cs
at 25 degrees C. The peak of 32 keV was clearly detected

4. Data Processing Board
We fabricated data processing FPGA board for telecommunication and control of sensor which made by Meisei
Electric Co., Ltd. The FPGA communicate with PC by
USB. It make house keeping data per 1 sec and multichannel photon (event by event) data per signal, andcontrol LSI, SRAM, clock, SRAM, ADC and HV with
Registered buﬀer.
We performed operation check of the FPGA andtook a spectrum using the FPGA from single
Ce:GAGG+APD and the VLSI unit. We will take multichannel data from multi unit arrays.
References
G. Sato, T. Kishishita, H. Ikeda, T. Sakumura, T.
Takahashi, IEEE Trans. Nucl. Sci, 58, 1370, 2011
Y. Yatsu et al., Proceeding of SPIE, Vol.9144, 2014
J. Kataoka et al., 2012, Journal of Instrumentation, 7,
6001

257

Advanced LIGO First Light: Multimessenger Astrophysics at the Birth ofGravitational-Wave Observatory
Leo Singer1
1

NASA Goddard Space Flight Center
E-mail(LS): leo.p.singer@nasa.gov
Abstract

Advanced LIGO’s direct observation of gravitational radiation from a binary black hole merger has sent
quakes through the physics and astronomy community. In a few short years, the search for gravitational
waves will complete its transformation from an experimental eﬀort into a new discipline of observational
astronomy as we rapidly build a sample of merging compact binaries. A particularly tantalizing goal is tocombine our new GW observatories with existing electromagnetic ones—uncovering the host environments
and formation channels of compact binaries, exposing the mechanism behind short GRBs, explaining thecosmic inventory of r-process elements, and even testing if stellar-mass black hole binaries are truly barrenof matter and magnetic ﬁelds. I will describe results from LIGO’s ﬁrst observing run with a special focus on
the electromagnetic follow-up program, which brings together real-time GW data analysis and 63 groups
who are searching for counterparts of LIGO so urces using ground- and space-based partner facilities
spanning gamma ray, x-ray, optical, infrared, and radio wavelengths, as well as neutrinos. I will discuss
the anticipated trajectory of the worldwide GW detector network from the standpoint of multimessenger
observations, and I will conclude with some ideas for the future.

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Gamma-ray
Bursts260

261

Gamma-ray bursts with MAXI and CALET
Satoshi Nakahira1 for the CALET collaboration
1

Japan Aerospace Exploration Agency (JAXA), 2-1-1 Sengen, Tsukuba, Ibaraki 305-8505
E-mail(SN): nakahira@crab.riken.jp
Abstract

MAXI and CALET are wide-ﬁeld observatories that have been operating on the same platform(JEM/EF of the ISS). They have overlapping ﬁelds of view and continuous energy ranges, and thus
GRB observations in a wide energy band can be achieved. From the start of the CALET observation in
October 2015 to December 2016, gamma ray burst monitor for CALET (CGBM) detected 49 GRBs, ﬁveof which were also detected with MAXI. In this paper, we will describe the gamma-ray burst observation
with CGBM and MAXI including simultaneously detected Gamma-ray bursts.
Key words: CALET, MAXI, GRB

1.

MAXI and CGBM observations of gamma-ray bursts

MAXI (Matsuoka et al. 2009) is a unique wide-ﬁeld monitor with its sensitivity in the soft X-ray band. It has been
observing all the sky since August 2009 to the present on
the exposed facility of Japan Experiment Module “Kibo”
on the international space station (ISS). MAXI has the
Solid State Camera (SSC) sensitive to 0.7–7 keV X-rays
and the Gas Slit Camera (GSC) sensitive to 2–30 keV
X-rays, which have rectangular ﬁelds of view of 90◦ and180◦ , respectively. MAXI has detected about 90 GRBs
in the past seven and a half years of observations. Half of
the GRBs detected by MAXI were found to be “MAXI
only bursts”, which were not detected by other instruments due to their dimness and soft spectra (Serino etal. 2014).
In October 2015 CALET (Torii et al. 2015) began observation on the same platform with MAXI. CALET
Gamma-ray burst monitor (CGBM; Yamaoka et al.
2013) has two types of scientiﬁc instruments: HXM(Hard X-ray Monitor) with a sensitivity in 7–1000 keV
and SGM (Soft Gamma Monitor) with a sensitivity in 50
keV–20 MeV. CGBM consists of two kinds of inorganic
scintillator crystals LaBr3 (Ce) (HXM1, 2) and BGO(SGM) with a PMT. Additionally CAL (CALorimeter; the main detector of CALET) has a sensitivity to
gamma-rays from 1 GeV to 10 TeV.
As summarized in Table 1, MAXI and CALET havecomplementary energy bands, and one-third of the
MAXI/GSC’s ﬁeld of view overlaps with all other instruments. Thus, GRB observations over an ultra-wide
energy band can be achieved.

Table 1. Properties of MAXI and CALET instruments.

MAXI

CALET

energy bands

0.7–7 keV (SSC)
2–30 keV (GSC)

7–1000 keV (HXM)
100keV–20 MeV (SGM)

instantaneousﬁelds of view

1% (SSC)

20% (HXM, CAL)

2% (GSC)

60% (SGM)

1 GeV–10 TeV (CAL)

2. Status of CGBM
In this section, we describe the performance of the
CGBM detectors, the status of calibration, data analysis and the operations on orbit.
2.1. performance
CGBM unit test was conducted in January 2014. We
performed a two-dimensional scan measurement of transmission eﬃciency and gain of the HXM entrance windowusing an X-ray generator and a radioisotope. Energy linearity of HXM and SGM direction dependence was alsomeasured. There was an unexpected asymmetric reduction of the transmission eﬃciency in the low energy bandof HXM 1 (Fig. 1).
Detection eﬃciency and gain at the higher energies
were almost uniform. Therefore, it is considered thatthe degradation of the eﬃciency occurs at the surfaceof the crystal. Energy dependency of the measured eﬃciency curve can be represented by absorption feature of
LaBr3 (Ce) itself. Possibly it is caused by a degradation

262

60

SGM
40

HXM1
HXM2

0

20

Effective Area (cm2)80

Total Efficiency

10100
Energy (keV)

1000

104

Fig. 2. Total eﬃciency of HXM1, 2 and SGM of CGBM (on-axis).
Fig. 1. HXM1 count rate map of 14.4 keV X-ray pencil bean taken
in ground test at January 2014.

• background integration time of 16s
of scintillating light emission eﬃciency due to hydrationof LaBr3 (Ce) in several μm of the surface by small air
leaks.
As for SGM, no major problem was found. Its lower
energy threshold on the orbit was ∼30 keV.
2.2. operation
CGBM regularly acquires two kinds of periodic data:• PH data: 1/8-second intervals with 512 channels• TH data: 4-second intervals with 8 channels
Also, when a GRB is triggered on the orbit,• capture of photon event data with 62.5μs time accuracy and 8192 channels• capture of an ASC (Advanced Stellar Compass) image with an exposure time of 0.5 second• change of CAL observation mode from HE mode(>10 GeV) to LE mode (>1 GeV)
will be executed.
The high voltage operation on orbit is controlled bya schedule command uploaded almost daily. Since theinclination angle of the ISS is as high as 51.6 degrees, toavoid false triggers and damage due to increase/variationof the charged particle rate, the trigger is turned oﬀ in
the high latitude region and the HV value is set to 0 V
subsequently. At the beginning of the CGBM operation,
we started with ∼45% of observation time referring to
the count rate map of the RBM (Radiation Belt Monitor)
equipped on MAXI. Now it is operating at observationeﬃciency of about 60%. The operating HV value is set
to 700.7 V, 793.8 V and 867.3 V for HXM 1, HXM 2 andSGM, respectively (gains of the ground measurements
can be approximated with these values).
Also, the trigger conditions are set as follows.

• foreground integration time of 1/4s, 1/2s, 1s, 4s• trigger energy range and signiﬁcance:– HXM 10–50 keV with 8.5σ
– SGM 40–230 keV with 7σ
2.3. data analysis
We are making eﬀorts to develop CGBM data in the Fits
format to meet the speciﬁcation of HEASoft1 provided
by HEASARC2 in order to simplify the future data publication.
The energy response function used for energy spectral analysis is created by a response generator using the Geant4 (Allison 2006) simulator. We modeled
CGBM detectors in detail and roughly reproduced whole
CALET geometry to reproduce back-scattered photons
from those structures. In order to take into account theeﬀects of degradation shown in section 2.1. in the calculation, we modeled the surface and the depth map of
the degraded part, which are used as calibration data.
Then this model is implemented in the response generator. The eﬃciency curve obtained from CGBM’s response matrix is shown in Figure 2.
We note that the response function still needs more
improvement because 1) optimization of HXM eﬃciency
by on-orbit data is not done yet and 2) eﬀects of ISS
structure including solar panels or robotic arms are notfully taken into account in the response calculation.

*1 http://heasarc.nasa.gov/lheasoft/*2 NASA’s High Energy Astrophysics Science Archive ResearchCenter263

3. MAXI and CGBM detected GRBs
In this section, we summarize the statistics of gamma-ray
bursts detected by MAXI and/or CGBM. Figure 3 shows
the angle distribution at the trigger time. In CGBM, 49
GRBs have been detected since the operation started, 40
of which are long, 9 are short, and 37 are localized byother satellites.
MAXI detected 15 GRBs during the same period. As
summarized in Table 2, ﬁve of them were also detectedwith CGBM, including one that came through the leadshield of the MAXI GSC.
Five of the CGBM-undetected GRBs were soft anddim, and the remaining one (GRB160705A) was a brightand soft burst, which emitted mainly below the sensitiveenergy bands of HXM. GRB160705A was found close to
the Galactic plane (b∼5◦ ), and it may be another kindof phenomenon such as an X-ray burst.

rate below 100 keV in HXM2 is lower compared to that of
HXM1, despite the greater eﬀective area of HXM2. An
increase in the plastic scintillator count rate in the CAL
Charge Detector (CHD) was also detected and reported
by Nakahira et al. (2016).

Table 2. Summary of CGBM detection of MAXI bursts.

References
Matsuoka, M. et al. 2009, PASJ, 61, 999
Serino, M. et al. 2014, PASJ, 66, 5, 87
Torii, S. for the CALET collaboration, Proc. 34th ICRC,
581 (2015)
Yamaoka, M. et al. 2013, proceeding of 7th Huntsville
Gamma-Ray Burst Symposium, C1304143
Allison, J. et al., 2006, IEEE Transactions on Nuclear
Science, 53,1, 270
Nakahira, S. et al. 2016, GCN, 19617

MAXI GRBs
detectedby CGBMnot detectedby CGBMinside of the MAXI’s FoV
outside of the MAXI’s FoVCGBM HV oﬀ
blocked by ISS structures
soft and dimbright and soft

15
4
1
2
2
5
1

3.1. Simultaneously detected events
We brieﬂy describe the characteristics of simultaneously
detected GRBs.
In GRB160101A and GRB160814A, the same peaks
were triggered simultaneously with MAXI and CGBM.
For GRB160101A, as shown in Figure 4(a), the pulse
width diﬀers depending on energy bands (The width isgreater at lower energies). As shown in Figure 4(d), thejoint spectrum is well reproduced with Band function
with α=−0.98±0.05, Ep =260.7+90−64 keV and a 2–1000
keV energy ﬂux of 1.69×10−6 erg cm−2 s−2 (β cannot
be constrained).
In the case of GRB160107A and GRB160509A, separated peaks were triggered. MAXI triggered the preburst emission of GRB160107A 45 seconds before themain peak detected by CGBM (Figure 4(b)), and caught
a bright X-ray ﬂare of GRB160509A during an earlyphase of the afterglow, 320 seconds after the main peakdetected with CGBM (Figure 4(c)).
GRB 160625A came from outside the MAXI ﬁeld ofview through the GSC lead shield. The spectra measured in CGBM are shown in Figure 4(e), corrected for
the HXM1, HXM2, and SGM eﬀective areas. However,
obscuration by the ISS solar panels is evident; the count

4. summary
Since October 2015, CALET/CGBM has detected 49
gamma-ray bursts, and ﬁve of them were also detectedby MAXI. Simultaneously observed bursts indicate different properties depending on the energy bands. The
MAXI and CGBM broadband observation is thought tobe useful for understanding gamma-ray bursts from different aspects.

264

180°

160927A

160821A

151006A
160623A
SGR1935+2154

90°160227A
160607A
160102A

SGR1935+2154
160118A
160104B

160223B

160509A

151205C 160509A

161004B

151112A

151216A

0°90

151107B
160625B160910AA
160408AA
160720A

60

160824B
160726A

160422A

160316A

30

160815A
160530B

161109A

160709A
161106A

160802A

151212B

151231A
A1160107A
160106A
151227B
160814A
151212A
151225A 151210B
1160705A
160419A
160101A

160206A
161123A

161010A

270°

120

front

Fig. 3. Incident angle distribution of GRBs detected by MAXI and CGBM. The red circles indicate the MAXI GRBs, the green crosses show
those of CGBM. The hatched regions in cyan and magenta are the ﬁelds of view of the GSC and SSC, respectively. The gray region shows
ISS structure viewing from the CGBM position.
(b) GRB160107A

50

í80

0
0

80

160

240 keV
HXM1+2 7-100

100

0

10

20

50

í80

30

SGM100-1000 keV

0
í10

0
10
20
TIME since 2016/01/01T00:43:53

0

80

160

120
counts /0.125sec

í10

240

SGM100-1000 keV

60
0
í80

30

0
80
160
TIME since 2016/01/07T22:20:42

240

SSC 0.7-7 keV

0.3
0.0
40

0

100

200

300

400

0

100

200

300

400

0

100

200

300

400

0

100

500

GSC2-20 keV 5.0

20
0
40

2.5
500

Eff.Area0

counts s−1 cm−2

counts /1.000 sec

10

0.0

0

25

í20

4

0.6

counts /2.0 sec

30
HXM1+27-100 keV

8

counts /0.125sec

20

GSC2-20 keV

20

counts /0.125sec

10

counts /0.125sec

0

effective Area

counts /0.250 sec
counts /0.125sec

counts /0.125sec

í10

100

0
í20
50

5.0
2.5

15
0
í20

(c) GRB160509A

30

GSC2-20 keV

Eff.Area(a) GRB160101A
30

0.0

HXM1+2 7-100 keV

20
0

200

500
SGM100-1000
keV

100
0

200
300
400
TIME since 2016/05/09T08:58:50

500(e) energy spectrum of GRB160101A
1(d) νFν spectrum of GRB160101A

0.01
10 −3
10 −4
10 −6

10 −5

Counts cm −2 s −1 keV −1

0.1

HXM1
HXM2
SGM

10

100

1000

10 4

Energy (keV)

Fig. 4. MAXI/GSC, CGBM/HXM and SGM light curves of (a)GRB160101A, (b)GRB160107A and (c)GRB160509A. Red lines in the
MAXI/GSC panels show the eﬀective areas. (d) νF ν spectrum of GRB160101A ﬁtted by the Band function and (e) energy spectraof GRB160625A corrected for the eﬀective areas.

265

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WKH6RIW*DPPDUD\'HWHFWRUDQG+DUG;UD\,PDJHURQERDUG+LWRPL
Koji Tanaka1 , Masanori Ohno1 , Yasushi Fukazawa1 , Tsunefumi Mizuno1 , Hiromitsu Takahashi1 , Takao
Kitaguti1 , Yasuyuki Tanaka1 , Jun’ichiro Katsuta1 , Takafumi Kawano1 , Hiroyuki Shirakawa1 , Ikumi
Edahiro1 , Sho Habata1 , Chiho Okada1 , Norie Ohashi1 , Tadayuki Takahashi2 , Motohide Kokubun2 , Shin
Watanabe2 , Goro Sato2 , Masayuki Ohta2 , Koichi Hagino2 , Rie Sato2 , Kunishiro Mori2 , Yusuke Uchida2 ,Miho Katsuragawa2 , Ryota Tomaru2 , Hiroki Yoneda2 , Kazuhiro Nakazawa3 , Shunsuke Torii3 , Soki
Sakurai3 , Makoto Sasano3 , Hidetoshi Nishida3 , Shogo Kobayashi3 , Hiroaki Murakami3 , Ko Ono3 , Yuichi
Kato3 , Kastuma Miyake3 , Yoshihiro Furuta3 , Yuki Murota3 , Hiroyasu Tajima4 , Masaomi Kinoshita4 ,Kazutaka Yamaoka5,6 , Katsuhiro Hayashi6 , Jun Kataoka7 , Tatsuhiko Saito7 , Masaki Yoshino7 , Takehiro
Mimura7 , Yoichi Yatsu8 , Shin’ya Saito9 , Hideki Uchiyama10, Takeshi Nakamori11, Shin’ichiro Takeda12 ,Toshio Nakano13, Takayuki Yuasa13, Yuto Ichinohe14, Daisuke Yonetoku15, and the Hitomi HXI/SGD team
1 Department of Physics, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan
2 Department of High Energy Astrophysics, Institute of Space and Astronomical Science (ISAS), Japan Aerospace

Exploration Agency (JAXA), 3-1-1 Yoshinodai, Chuo, Sagamihara 252-5210, Japan
3 Department of Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
4 Solar Terrestrial Environment Laboratory, Nagoya University, Nagoya, Aichi 464-8601, Japan
5 Institute for Space-Earth Environmental Research, Nagoya University,

Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8601, Japan
6 Division of Particle and Astrophysical Science, Graduate School of Science, Nagoya University,

Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8602, Japan
7 Research Institute for Science and Engineering, Waseda University, Shinjuku-ku, Tokyo, 169-8050, Japan
8 Department of Physics, Tokyo Institute for Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8550, Japan
9 Department of Physics, Rikkyo University, Toshima-ku, Tokyo, 171-8501, Japan
10 Science Education, Faculty of Education, Shizuoka University, Suruga-ku, Shizuoka, 422-8529, Japan
11 Department of Physics, Faculty of Science, Yamagata University, Kojirakawa, Yamagata 990-8560, Japan
12 Advanced Medical Instrumentation Unit, Okinawa Institute of Science and Technology ]Graduate University, Onna, Okinawa 904-0495, Japan
13 Cosmic Radiation Laboratory, Institute of Physical and Chemical Research (RIKEN) Wako, Saitama 351-0198, Japan
14 Department of Physics, Tokyo Metropolitan University, Hachioji, Tokyo 192-0397, Japan
15 Department of Physics, Kanazawa University, Kanazawa, Ishikawa 920-1192, Japan

E-mail (MO): tanakak@astro.hiroshima-u.ac.jp

ABSTRACTSoft Gamma-ray Detector (SGD) and Hard X-ray Imager (HXI) onboard Hitomi are surrounded bylarge and thick Bi4Ge3O12(BGO) crystal scintillators to reduce background. These BGO scintillators act
as not only active shield but also all-sky monitor, especially for Gamma-ray Bursts (GRBs). Therefore,
we estimated the capability of localization of GRBs with SGD and HXI BGO active shield by using
Geant4 based Monte Carlo simulation. We developed simulator of SGD and HXI BGO active shieldincluding mass model of Hitomi and detector parameters which are obtained ground calibration test
using flight model sensors. We construct database of detection efficiency of GRBs for all 68 BGO units.
We established the localization method which is finding best fit incident angle by comparing observed
photon counts of all 68 units and calculated database. By applying this method to actually observed
GRB, GRB160324A by HXI/SGD BGO active shield, we successfully constrained incident photon
direction within 16 degrees uncertainty.
KEY WORDS: Gamma-ray Burst , Hitomi: HXI/SGD: Active Shield

268

1.

Introduction

The origin of short GRBs (their duration of lessthan about two seconds) is believed to be a mergerof compact objects such as neutron star and blackhole, and they are candidate of the source ofgravitational wave (GW). Short GRBs and GW
connection could reveal the origin of short GRB.
Thus, localization of GRBs are very important.
SGD and HXI BGO active shield have large
geometrical area, and they are expected to
localize GRBs utilizing the difference of thedetection efficiency
between total 68 BGO
detectors which face different sky direction.
The aim of this study is to establish GRB
localization methods utilizing SGD and HXI BGO
active shields by using Geant4 based Monte Carlosimulation.
2.

Overview of GRB Localization Method

2.1. Development of Simulator
We have to develop a Monte Carlo simulator
whitch can reproduce the complicate detector
responses of BGO detectors, implementing
actually measured detector parameters such asenergy resolution and so on. Detailed detector
mass model is also constructed. Here we utilized
Geant4 toolkit (ver. 10.2.p2).

2.3. GRB Localization
Here, the chi2 is defined as following andobserved GRB photon counts of all 68 BGO units
and constructed database are compared.
ሺ ‫୳ כ‬୬୧୲ ሺߠǡ ߮ሻ െ ୳୬୧୲ ሻଶ߯ଶ ൌ ෍ ቆቇଶ
୳୬୧୲where norm, incident zenith and azimuthal angle,
ߠ and ߮ are the fitting parameter. modelunit anddataunit are the observed and predicted count rate,
respectively. The best fit incident angle can be
estimated by minimizing this ߯ ଶ value.

3.

Localization of GRB160324A

We applied this method to GRB160324A observedby Hitomi and tried to localize.
Figure 2 shows the result of localization.We
found that the incident zenith angle ߮ of
GRB160324A is 46.9ι(+2.7,-2.0), and azimuthal
angle ߠ is 42.5ι(+5.9,-2.6). This result is roughly
consistent with IPN localization (private
communication).
We estimated systematic error of this methodby 100 times realizations, and it is found to beabout 10 degrees.

2.2. Detection Efficiency Database
We constructed detection efficiency database byscanning simulated GRBs all 68 BGO units.
Figure 1 shows an example of detection efficiency
database (only SGD1 units), and found from theresult that the units face different direction havedifferent detection efficiency pattern.
Figure 2. Localization result for GRB 160324A. It
shows the chi2 distribution for incident azimuthaland zenith angle. The star mark shows pointed thebest localization angle.

4.

Figure 1. Incident angle dependence of detection
efficiency of 25 BGO units. Horizontal and verticalaxis of each panel represents incident azimuthaland zenith angle of GRB, respectively.

Conclusion

We established the localization method based on
simulated detection efficiency database, and it
works well for localization of actually observed
GRB 160324A. The obtained incident angle isroughly consistent with other results. Systematic
uncertainty of this localization can be estimatedto be about 10 degrees in this case.

269

Suzaku Wide-band All-sky Monitor view of soft gamma-ray transient
Masanori Ohno1 , Kazutaka Yamaoka2,3 , Makoto S. Tashiro4 , Kevin Hurley5 , Hans A. Krimm6,7 , Amy Y. Lien6,8 ,Norisuke Ohmori9 , Satoshi Sugita10 , Yuji Urata11 , Tetsuya Yasuda4 , Ryuuji Kinoshita9 , Yusuke Nishioka9 ,Yoshitaka Hanabata12 , Jun’ichi Enomoto4 , Takeshi Fujinuma4 , Yasushi Fukazawa1 , Wataru Iwakiri13 ,Takafumi Kawano1 , Motohide Kokubun14 , Kazuo Makishima15 , Shunsuke Matsuoka4 , Tsutomu Nagayoshi4 ,Yujin Nakagawa16 , Souhei Nakaya4 , Kazuhiro Nakazawa17,18 , Takanori Sakamoto19 , Tadayuki Takahashi14 ,Sawako Takeda4 , Yukikatsu Terada4 , Seiya Yabe4 , and Makoto Yamauchi91

Department of Physics, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan
2
Institute for Space-Earth Environmental Research, Nagoya University,
Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8601, Japan
3
Division of Particle and Astrophysical Science, Graduate School of Science,
Nagoya University, Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8602, Japan
4
Department of Physics, Saitama University, 255 Shimo-Okubo, Sakura-ku, Saitama-shi, Saitama 338-8570, Japan
5
Space Sciences Laboratory, University of California, 7 Gauss Way, Berkeley, CA 94720-7450, USA
6
Center for Research and Exploration in Space Science and Technology (CRESST)
and NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
7
Universities Space Research Association, 7178 Columbia Gateway Drive, Columbia, MD 21046, USA
8
Department of Physics, University of Maryland, Baltimore County, 1000 Hilltop Circle, Baltimore, MD 21250, USA
9
Department of Applied Physics, University of Miyazaki,
1-1 Gakuen kibanadai-nishi, Miyazaki-shi, Miyazaki 889-2192, Japan10
Department of Physics, Tokyo Institute for Technology, 2-12-1
Ookayama, Meguro-ku, Tokyo 152-8550, Japan
11
Institute of Astronomy, National Central University, Chung-Li 32054, Taiwan
12Institute for Cosmic-Ray Research, University of Tokyo, 5-1-5 Kashiwa-no-ha, Kashiwa-shi, Chiba 277-8582, Japan
13
MAXI Team, The Institute of Physics and Chemical Research (RIKEN), 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
14
Department of High Energy Astrophysics, Institute of Space and Astronomical Science (ISAS),
Japan Aerospace Exploration Agency (JAXA), 3-1-1 Yoshinodai, Chuo, Sagamihara 252-5210, Japan
15
RIKEN Nishina Center, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
16
Science-satellite Operation and Data Archive Unit, Institute of Space and Astronomical
Science (ISAS), Japan Aerospace Exploration Agency (JAXA), 3-1-1 Yoshinodai, Chuo-ku,
Sagamihara-shi, Kanagawa 252-5210, Japan17
Department of Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan18
Research Center for the Early Universe, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
19
Department of Physics and Mathematics, Aoyama Gakuin University,
5-1-1 Fuchinobe, Sagamihara, Kanagwa 252-5258, Japan
E-mail(MO): ohno@astro.hiroshima-u.ac.jp
Abstract
In this paper, we review results of soft gamma-ray transients observed by the Wide-band All-sky
Monitor (WAM) onboard Suzaku satellite. The WAM had been operated for about 10 years from 2005 to
2015, and observed various types of transients and bright gamma-ray sources, such as gamma-ray bursts(GRBs), solar ﬂares, soft gamma repeaters (SGRs), gamma-ray emission from Crab nebula ..etc. Thanks
to its very large eﬀective area up to soft gamma-ray energy band (∼400 cm2 even at 1 MeV), we can detect
soft gamma-ray photons from many objects with good statistics. Especially, detected number of GRBs is
more than 1400, which is comparable detection rate to other GRB instruments. With those large numberof samples, we can discuss many kinds of science topics. For example, the duration distribution of observed
GRBs, timing and spectral properties for GRBs, weak GRBs with high redshifts, time-resolved pulses forlong GRBs, spectral properties of solar ﬂares in hard X-ray energy band, and emission mechanism of brightbursts from SGRs. In addition to transient observations, the WAM has the capability to monitor the bright
gamma-ray sources utilizing earth occultation technique. We successfully monitored some sources such as
Crab, Cyg X-1 from 150 to 500 keV energy band and we can trace the broad-band spectral variability.
Key words: Suzaku:HXD-WAM — all-sky observations

270

The Hard X-ray Detector (HXD(Takahashi et al.
2007)), which is one of two narrow-ﬁeld instruments onboard Suzaku satellite, is surrounded by very large 20
BGO crystals. The primary role of these BGO crystals
is to serve main detector as an active shield. In addition,
their large geometrical, eﬀective area and wide ﬁeld ofview enable us to observe all-sky covering very wide energy band (Wide-band All-sky Monitor; WAM(Yamaoka
et al. 2017)). The WAM has the largest eﬀective area
above 300 keV among any other all-sky instruments asshown in Fig 1.

Fig. 1. Comparison of eﬀective area with other all-sky instruments

The WAM has been operated for ∼ 10 years from August 2005 to May 2015. There were 2935 onboard triggers throughout the mission, consisting of 856 GRBs conﬁrmed by other instruments, 309 possible GRBs judged
from the light curve, 48 SGRs, 81 solar ﬂares and so on.
Total detected number of GRBs, including un-triggered
events is more than 1400 and an average detection rateis ∼140 GRBs per year. Utilizing such large numberof GRB samples and thanks to the large eﬀective areaof the WAM, detail studies for spectral, temporal, andmulti-wavelength properties of GRBs have been reported
through more than 30 papers.
Spectral parameters of 193 GRB samples reported
in the GCN circulars was compared by Yamaoka etal. (2017) and found that the WAM Epeak distribution
has a mean around 700 keV, much larger than that of
Swift/BAT and Fermi/GBM. This is due to instrumental
eﬀect that the low energy coverage of the WAM is slightly
higher than other instruments. Krimm et al. (2009)
conﬁrmed that BAT and WAM joint analysis shows themean Epeak value of around 250 keV, consistent withother results. The Epeak evolution for both individual
and inside pulses are studied by Sugita (2009) and Ohno
et al. (2009) and they found the clear correlation between time-resolved Epeak and isotropic luminosity Liso .
Based on synchrotron shock scenario, this suggests constant emission radius. The WAM also observed high-z
GRBs, GRB 050904(z=6.3), GRB 120521C(z∼6.0), and
GRB 130606A(z∼5.9), and found that they also follow
the same Epeak − Eiso correlation as normal GRBs.
One of the most important temporal property to dis-cuss its origin is the duration distribution. The detail study using WAM sample has been reported byOhmori et al. (2016) and conﬁrmed clear bimodal,
short/long GRB separation, which is similar to that of
INTEGRAL/SPI-ACS and BATSE. The origin of differences from other instruments such as Swift/BAT isalso discussed. Tashiro et al. (2012) studied the temporal decay properties of GRBs and revealed that theenergy dependence of the decay constant could be explained by the synchrotron emission model. Combining
to other wavelength information of afterglow emission
allows us to estimate jet-collimation corrected enegetics.
Urata et al. (2009, 2012) examined the correlation between jet-collimation corrected energy Eγ and Epeak andfound an outlier for GRB 071010B. While, Fermi-LAT
detected GRBs are found to be consistent with classicallong GRBs despite of their energetic GeV photon emitting property. Larger sample could reveal a detail of suchoutlier and relation to GeV emission property.
The WAM also observed many non-GRB transients.
Yasuda et al. (2015) reported that the detection of MeV
photons from outburst of SGR J1550-5418 on January
22th 2009. The spectrum is represented by black bodyplus bremsstrahlung by optically thin plasma or cut-oﬀ
power-law model. Such high energy photon could reveal
environment properties of SGRs. The overall properties
of WAM solar ﬂares is discussed in catalog papers (Endoet al. 2010, Yabe et al. in this volume). In addition, detail spectral properties is also reported for joint analysis
with RHESSI by Ishikawa et al. (2013), and accelerated
electron properties are discussed.
The WAM also has a capability to monitor bright
gamma-ray sources utilizing earth occultation technique.
Some bright gamma-ray sources such as Crab and Cyg
X-1 are successfully monitored for more than 7 years,
and time-averaged Crab spectrum is detected up to 700
keV thanks to the large eﬀective area of the WAM. The
spectrum is consistent with other energy band obtainedby the main detector of HXD. The long-term light curveof the Crab is also consistent with other instruments suchas Fermi/GBM. More detail result of the earth occultation observations will be published.
References
Takahashi, T. et al. 2007 PASJ, 59, S35
Krimm, H. A. et al. 2009 ApJ, 704, 1405
Sugita, S. 2009, phD thesis
Ohno, M. et al. 2009 PASJ, 61, 201
Endo, A. et al. 2010 PASJ, 62, 1341
Ishikwa, S. et al. 2013 ApJ, 765, 143
Urata, Y. et al. 2012 ApJ, 748, L4
Urata, Y. et al. 2009 ApJ, 706, L183
Yasuda, T. et al. 2015 PASJ, 67, 41
Tashiro, M. S. et al. 2012 PASJ, 64, 26
Ohmori, N. et al. 2016 PASJ, 68, 30
Yamaoka, K. et al. 2017 PASJ, submitted

271

Energy response function of CALET Gamma ray Burst Monitor
Y. Yamada,1 T. Sakamoto,1 A. Yoshida,1 for the CALET collaboration
1

Department of science and engineering, Aoyama Gakuin University
5-10-1 Fuchinobe Chuo-ku, Sagamihara-shi, Kanagawa-ken, 252-5258, Japan
E-mail(FS): yyamada@phys.aoyama.ac.jp
Abstract
CALorimetric Electron Telescope (CALET) was successfully launched and attached to the exposed
facility of the Japanese Experimental Module (JEM) called Kibo at the International Space Station (ISS)on August, 2015. CALET started its regular operation on Octorber 2015. CALET has two scientiﬁc
instruments: CALorimetor (CAL) and CALET Gamma-ray Burst Monitor (CGBM). CGBM is speciﬁcallydesigned to observe gamma-ray bursts (GRBs).
We calibrated the energy response function of the CGBM using the ground and the ﬂight data. The
well calibrated energy response function is crucial for the spectral analysis of GRBs. We constructed themass model of CALET using the GEANT4 simulation package. This mass model is used to construct theenergy response function of CGBM. In this paper, we present the current status of the energy response
function of the CGBM.
Key words: CALET
1. CALET Gamma-Ray Burst Monitor
CALorimetric Electron Telescope (CALET) was successfully launched and attached to the exposed facility of the
JEM called Kibo at the International Space Station (ISS)on August, 2015. (Torii et al. 2015) CALET started its
regular operation on Octorber 2015. CALET has two
scientiﬁc instruments: CALorimetor (CAL) and CALET
Gamma-ray Burst Monitor (CGBM). CGBM is specifically designed to observe gamma-ray bursts (GRBs).
CGBM covers a wide range of energy by using two different types of scintillation detectors: The Hard X-ray
Monitor (HXM), based on LaBr3 (Ce), covers the energy
range 7 keV ∼ 1 MeV. The Soft Gamma-ray Monitor(SGM), based on BGO, covers the range 50 keV ∼ 20
MeV. (Yamaoka et al. 2013)
2. CGBM simulator and Full CGBM simulator
We use the GEANT4 (GEometry ANd Tracking 4) simulation package to construct the mass model of CGBM.
The physical processes included in this simulator are thephotoelectric eﬀect, Compton scattering, electron-pair
production and Rayleigh scattering.
2.1. CGBM simulator
Individual HXM and SGM simulators are built at ﬁrst.
We compared the simulator to the ﬂight model. We
checked the mass, the size, the material of the simulator
to the actual design of the ﬂight model and they agreed
well with the ﬂight model.

Fig. 1. CALET simulator as viewed from the top.

Fig. 2. CALET simulator as viewed from the side.

2.2. Full CGBM simulator
Next, we built the full CGBM simulator which includes
the basic structures of CALET. In order to construct
the energy response function, the eﬀect of scattering by
the surrounding materials is needed to be simulated accurately. Therefore, it is necessary to build the entire
structure of CALET into the simulation, which we label
the full CGBM simulator.
We incorporated the entire structure of CALET into

272

104

100

Blue: Ground data
Red: Simulation data90

+ 0 deg+ 10 deg+ 30 deg+ 60 deg+ 90 deg

80

Effective Area [cm2]Counts

103

102

10

70
60
50
40
30

1

20
10−1
0

200

400

600

800

1000

1200

1400

1600

1800

10

2000

Energy [keV]
1

3

102

10

104

10

Energy [keV]

105

Blue: Ground data
Red: Simulation data

104

120

Effective Area [cm2]Counts

102

10

1

10−1
0

+ 0 deg+ 30 deg+ 60 deg+ 90 deg

100

103

80

60

40

200

400

600

800

1000

1200

1400

20

Energy [keV]
102

3

10

104

Energy [keV]

Fig. 3. The simulation spectrum (22 Na) above: The CGBM simulation bottom: The full CGBM simulation
Fig. 4. Eﬀective area of HXM (above) and SGM (bottom) of the full
CGBM simulator

the full CGBM simulator, including CGBM, CAL, ASC,
MDC, E-Box, and HV-Box. We checked the mass, thesize, the material of the simulator and their positions ofincorporated substances. They agreed well to the ﬂight
model. The ﬁgure 1, 2 are the full CGBM simulator.
3.

Comparison between ground calibration data and simulation data
We performed ground calibration tests to obtain the datafor constructing the energy response function. We used
the following radiation sources: 22 Na, 57 Co, 88 Y, 109 Cd,
133
Ba and 137 Cs. We measured the energy linearity, theangle response and the uniformity of the crystal surface.
3.1. Veriﬁcation of the CGBM simulator
The radiation source, 22 Na, was located in front of the
CGBM for the simulation. The position of the radiation
source is the same as the setup of the ground test. We
replicate the setup of the ground calibration set up in
the simulation. We simulated and made a comparison
between the ground data and the simulation data. The
simulation data and the ground calibration data agree
quite well (ﬁg 3).
3.2. Veriﬁcation of the full CGBM simulator
The radiation source (22 Na and 109 Cd) were located on
top of the CALET. The heights of radiation source are
0.3 m, 0.5 m and 0.86 m from the top part of the CGBM.

We replicate the setup of the ground calibration test in
the simulation. We compared the measured data andsimulated data. The low energy part shows a mismatch
between the simulation and the calibration data (ﬁg 3).
We are currently investigating the reason for this discrepancy. One possible reason is a lack of Compton scattering eﬀect by surrounding materials in the simulator.
4. Eﬀective area
We measured the eﬀective area at several incident angles
for the SGM and HXM using the full CGBM simulator.
The eﬀective area was calculated by changing the incident angles (0, 30, 60 and 90 deg). Both the CGBM
simulation and the full CGBM simulation have the maximum eﬀective area at 30 degrees, and when the angle
is 90 degrees, eﬀective area is the smallest. The eﬀective area matches when the angle is 0 degrees, but when
we changed the incident angle, the eﬀective area of thefull CGBM simulator becomes smaller than the CGBM
simulator (ﬁg 4).
References
Torii, S et al. PoS (ICRC2015) 851
Yamaoka, K et al. 2013, 7th Hunsville Gamma-Ray
Burst Symposium, GRB 2013: paper 41 in eConf
Proceedings C1304143

273

Overview and ﬁrst year observation of CALET Gamma ray Burst Monitor
Yuta Kawakubo,1 Atsumasa Yoshida,1 Takanori Sakamoto,1 for the CALET collaboration
1

Department of Physics and Mathematics, Aoyama Gakuin University
5-10-1 Fuchinobe, Chuo, Sagamihara, Kanagawa 252-5258, Japan
E-mail(YK):ykawakubo@phys.aoyama.ac.jp
Abstract

Calorimetric electron telescope (CALET) was launched by the 5th H-IIB/HTV on August 19, 2015.
CALET includes CALET Gamma ray Burst Monitor (CGBM) to observe GRB prompt emissions. CGBM
consists of two Hard X-ray Monitors (HXM) and Soft Gamma ray Monitor (SGM). CGBM is able toobserve prompt emissions of GRBs in the 7 keV - 20 MeV band (HXM; 7 keV - 1 MeV, SGM; 40 keV 20 MeV). CALET started its nominal observation on October 8, 2015. So far, CGBM detected 49 GRBs
which include 9 short GRBs and 40 long GRBs. The detection rate is about 4 GRBs per month. Four
GRBs are also detected by MAXI. The distribution of T90 is similar to GRBs detected by BATSE. The
mean of T90 measured by CGBM is 8.9 s.
Key words: GRB — CALET
1. Introduction
A gamma ray burst (GRB) is a mysterious explosionin the universe. A prompt emission, an explosive main
emission of GRBs, is very bright and variable. Although
the mechanism of a prompt emission is not well understood, a cogent scenario is the internal shock model.
That is collisions between each ejecta which have a different velocity in a relativistic jet cause the internal
shock (Rees and Meszaros, 1994 ). The shock accelerated electrons radiate X-rays or gamma-rays due to thesynchrotron radiation. The spectra of the prompt emissions show a broad energy distribution from keV to MeV
range. These spectra are well represented by a smoothly
jointed broken power-law, namely Band function (Bandet al. 1993 ). To study its wide band nature, spectral
observations of prompt emissions need a broad energy
coverage. Fermi/LAT detected very high energy photons above 100 MeV from GRBs and found out that theycame later than typical hard X-ray photons of prompt
emission (Abdo et al. 2009 ). The high energy photons
are thought to be caused by an external shock whicharises when the ejecta from the GRB collides with interstellar medium.
2. CALorimetric Electron Telescope (CALET)
A main purpose of CALorimetric Electron Telescope(CALET) mission is observations of cosmic electrons
and gamma-rays in the GeV - TeV energy band (Torii
et al. 2011 ). CALET also aims to observe high energy transients such as GRBs in X-ray and gamma-

Fig. 1. The image of CALET. CALET is equipped with one SGM andtwo HXMs.

ray range. CALET carries two scientiﬁc instruments
CALorimetor (CAL) and CALET Gamma ray Burst
Monitor (CGBM). CAL has a sensitivity for gamma-raysin the 1 GeV - 10 TeV band. CGBM utilizes two typeof scintillation detectors : Hard X-ray Monitor (HXM;
7 keV - 1 MeV) and Soft Gamma ray Monitor (SGM;
40 keV - 20 MeV). All sensors of CGBM have two ampliﬁers with diﬀerent gain (high / low) to observe GRB
radiations in the broad energy range.
Table 1. Speciﬁcation of HXM and SGM
Number of detector
Detector
Diameter [mm]Thickness [mm]
Energy range [keV]
Field of view

HXM
2
LaBr3 (Ce)61
12.7
7-1000
∼3 sr

SGM
1
BGO
102
76
100-20000
∼8 sr

274

Table 2. Data of CGBM

Data
Time History (TH)Pulse Height (PH)Event data (EVT)

Channels (High)
4
102
4096

Channels (Low)
4
410
4096

Time resolution
0.125s4s
62.5μsCALET was launched by the 5th H-IIB/HTV on August 19, 2015 and attached on the JEM module of theinternational space station (ISS) on August 25. The initial checkout was performed on September 22. The 72
hour running test was carried out from October 5 to 8.
During this test, CGBM detected GRB 151006A which
was the ﬁrst GRB detection by CGBM. CGBM started
the nominal observation on October 8. Since CGBM
high voltages are turned oﬀ at the SAA region and highlatitude regions to protect PMT and to reduce false trigger alerts, the observation eﬃciency is about 60%.
4. Observation
CGBM detected 49 GRBs, 40 long bursts and 9 shortbursts, during October 2015 to November 2016. The
detection rate of GRBs is ∼4 GRBs per month. The
most of the CGBM GRBs were also detected byother instruments (e.g. Fermi/GBM, INTEGRAL/ACS,
KONUS/Wind, Swift/BAT, MAXI/GSC). Especially, 4
GRBs were also detected by MAXI/GSC. The T90 durations of GRBs detected by CGBM were measured by
SGM in the 40 - 1000 keV band. We classiﬁed GRBs
based on the T90 smaller than 2 s as short bursts. Fig. 2
shows the T90 distributions of GRBs detected by CGBM.
The mean of T90 of the CGBM GRBs is 8.9 s a bit
smaller than that of 9.3 s by BATSE (see Fig. 2). Since
the T90 is measured at diﬀerent energy bands by diﬀerent instruments, the measured T90 could be aﬀected by
the energy dependence of GRB pulses (Fenimore et al.
1995 ) and the sensitivity of instruments.
CALET joins to the LIGO/Virgo electro-magnetic follow up program. When GW 150914 occurred, CGBM

!"#$#" 

3. Data collection and In-orbit operation
CGBM regularly collects two kinds of monitor data. One
is Pulse Height (PH) data. The other is Time History (TH) data. CGBM has the on-board trigger system which calculates a signal-to-noise ratio every 0.25
s. If the signal-to-noise ratio exceeds a trigger threshold,
CGBM captures event data with a time resolution of 62.5
μs. We setup a pipeline to send an alert to Gamma-ray
Coordinates Network automatically. When the CGBM
on-board trigger happens, CAL decreases the lower energy thresholds for gamma ray from 10 GeV to 1 GeV.
Advanced Stellar Compass (ASC) which is a optical starcamera takes two images.


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Fig. 2. The distribution of T90 . The top panel shows distribution of
CGBM GRBs (mean is 8.9s). The middle high panel shows thatof BATSE GRBs (mean is 9.3s). The middle low panel shows thatof Fermi/GBM GRBs (mean is 13.4s). The bottom panel showsthat of Swift/BAT GRBs (mean is 24.4s).

was not in the normal operation. However, when GW
151226 occurred, CGBM was working. We checked TH
light curve data within a period of ± 60 s from the GW
trigger time. As the result, no signiﬁcant signal was seenin the CGBM data. We also estimated 7σ upper limit
for GW 151226 assuming a spectrum of a typical short
GRB (α= −0.58, Epeak = 355 keV; Ghirlanda et al. 2009). The 7σ upper limit of HXM is 1.0 ×10−6 cm−2 s−1(7 - 500 keV) assuming 1 s integration time. The 7 σ
upper limit of SGM is 1.8 ×10−6 cm−2 s−1 (50 - 1000keV) assuming 1 s integration time (Adriani et al. 2016).
5. Summary
CGBM is performing a nominal observation since October 8, 2015. The current observation eﬃciency is ∼60%. CGBM detected 49 GRBs which consist of 40 long
GRBs and 9 short. The rate of short GRBs is ∼18 %.
The detection rate of GRB is 4 GRBs a month. 4 GRBs
were also detected by MAXI/GSC. The distribution of
the T90 obtained by CGBM is similar to that by BATSE.
CALET did not ﬁnd a signiﬁcant electro magnetic counterpart of GW151226 (Adriani et al. 2016 ).
References
Abdo, A. A. et al. 2009 Science 323, 1688
Adriani, O. et al. 2016 ApJ 829, 20
Band, D. L. et al., 1993 ApJ, 413, 281
Fenimore, E. E. et al. 1995 ApJ, 448, 101
Ghirlanda, G. et al., 2009 AA, 496, 585
Rees, M. J. and Meszaros, P. 1994 ApJ 430, 93
Torii, S. et al., 2011 NIM A, 2011, 630, 55

275

Extremely Soft X-ray Flash as the indicator of oﬀ-axis orphan GRB
afterglow
Yuji Urata,1 Kuiyun Huang,2 , Ryo Yamazaki,3 Takanori Sakamoto3
1

2

Institute of Astronomy, National Central University, Chung-Li 32054, Taiwan
Department of Mathematics and Science, National Taiwan Normal University, Lin-kou, New Taipei 2449, Taiwan
3
Department of Physics and Mathematics, College of Science and Engineering, Aoyama Gakuin University,
5-10-1 Fuchinobe, Chuo-ku, Sagamihara-shi, Kanagawa 252-5258, Japan
E-mail(FS): urata@astro.ncu.edu.tw
Abstract
Toward uniﬁcation of GRBs and X-ray Flashes (XRFs), we veriﬁed the oﬀ-axis jet model and examineda discovery of oﬀ-axis orphan GRB afterglows. The XRF sample was selected on the basis of the following
three factors: (1) a constraint on the lower peak energy of the prompt spectrum, (2) redshift measurements,and (3) multicolor observations of an earlier (or brightening) phase. XRF020903 was the only sample
selected basis of these criteria. An oﬀ-axis jet model with a large observing angle (0.21 rad, which is
twice the jet opening half-angle), can naturally describe the achromatic brightening and the prompt X-ray
spectral properties. This result indicates the existence of oﬀ-axis orphan GRB afterglow light curves.
Events with a larger viewing angle (>2 times of the jet opening angle) could be discovered using Subaru
Hyper-Suprime-Cam (HSC), LSST, and SKA.
Key words: Gamma-Ray Bursts — Oﬀ-axis orphan GRB afterglow — X-ray Flash1.

Introduction

GRBs are believed to be caused by jet collimation andultra-relativistic motion. As identical to that of AGN
classiﬁcation, a uniﬁed picture of GRBs and massive stellar explosions is expected. Along to the observing angleof GRB jet, four astronomical phenomena should exist;(A) classical GRBs, (B) X-ray ﬂash (XRFs), (C) oﬀ-axis
orphan GRB afterglows, and (D) supernova. However,no direct, unambiguous, observational evidence exists.
Toward establishment of the uniﬁcation model, (1) demostration of oﬀ-axis origin of XRFs and (2) discovery
of oﬀ-axis orphan GRB afterglows are required.
XRFs exhibit that X-ray temporal and spectral properties in the prompt phase do not show any diﬀerences
relative to those of GRBs, except for the considerably
lower energy values of the peak of the νFν spectrum inthe observer’s frame. The majority of HETE-2 samples(nine out of 16 XRFs) show a low energy of spectral peakobs
energy Epeak< 20 keV. The number of XRFs detected
by HETE-2 was comparable and relatively larger thanthat of GRBs indicating that XRFs represent a large
portion of the entire GRB population. To explain the
aforementioned prompt observational properties, threemodels have been proposed for XRFs: a high redshift
origin; the oﬀ-axis jet model, which is equivalent to theuniﬁcation scenario of AGN galaxies; and intrinsic properties (e.g., a subenergetic or ineﬃcient ﬁreball), which
may also produce on-axis orphan afterglows.
Oﬀ-axis orphan afterglows occur as a natural prediction of GRB jets and ultra-relativistic motion. However,no direct, unambiguous, observational evidence exists.
We therefore verify the oﬀ-axis jet model and to provide feedback to ongoing and planned optical time domain surveys of Subaru Hyper-Suprime-Cam (HSC) and
LSST (Urata et al. 2015).
2. XRF020903
We considered possible XRFs for our study and quickly
realized XRF 020903 was the only event that has (1)either a measurement or an upper bound on the peaksrc
energy of the prompt spectrum Epeak, (2) a measured
redshift, and (3) multicolor afterglow observations adequately cover the early afterglow phase when achromaticbrightening of the afterglow might occur.This was one of
the XRFs detected by HETE-2 and the ﬁrst events forwhich an optical afterglow was detected and a spectroscopic redshift (z = 0.251) was determined. The prompt
emission had the lowest intrinsic spectral peak energy
src
Epeak
of 3.3+1.8−1.0 keV among all the XRF samples. We
collected data for the XRF 020903 afterglow by using

276

0.1

0.01

10í3
0.1

1

10

100

Days after the burst

Fig. 1. Multi-frequency light curves of XRF020903 afterglow. The
dotted lines show the oﬀ-axis jet model functions.

Subaru archive data and photometric results obtained
from literature. To perform accurate optical photometry by removing contamination from the host galaxy, we
also employed Pan-STARRS1 data.
Figure 1 shows the optical and radio (8.5 GHz) light
curves. The optical (Rc and Ic) light curves show aclear rapid rebrightening between 0.7 and 0.9 days. The
equivalent rising power law index α (Fν ∝ tα ) are ∼ 2.6
in the Rc and ∼ 2.3 in the Ic band. There is anotherlate-phase rebrightening peaking at ∼ 20 days in the Rband light curve, which was interpreted as the associatedsupernova component. The SED at 0.7 and 0.9 day iswell ﬁtted by the power-law function as f (ν) ∝ ν −β . We
obtained the β at 0.695 day as 1.48 ± 0.06. Similarly,
we obtained a β as 1.43 ± 0.08 at 0.899 day, which isconsistent with the value for 0.695 day. These results
imply that the rebrightening is achromatic.
To perform light curves and SEDs modeling for
XRF020903, we employed the boxﬁt code. Figure 1
shows the best-ﬁt model functions that describe theachromatic rebrightening, optical temporal evolution,and radio brightness. There are two notable features.
The ﬁrst is the jet opening half angle θjet of 0.1 rad,which is consistent with those of classical hard GRBs.
The second is the large observing angle (θobs ∼ 0.21 rad),which corresponds to θobs ∼ 2θjet .
src
andWe veriﬁed the observed small values of Epeak
Eiso by adopting a simple model with a top-hat proﬁleof the prompt emission of a relativistic jet. The paramesrc
ters Epeak(θobs ) and Eiso (θobs ) were analytically derived
as functions of θobs , θjet , and γ. By ﬁxing θobs and θjetsrc
as 0.21 and 0.1 rad, respectively, we evaluated γ, Epeak,src
and Eiso observed from the on-axis of the jet (Epeak (0)and Eiso (0)). These expected values were consistent withthose of classical hard GRBs, and the observed small valsrc
ues of Epeak
and Eiso of XRF 020903 could be naturally
explained by the oﬀ-axis jet model.

3. Prospect with Optical Surveys
We expected oﬀ-axis orphan GRB afterglow light curves
at z = 0.5 along with three viewing angles on the basisof the XRF020903 result (Figure 2 top). To detect theselight curves, especially afterglows with a larger viewing
angle (θobs > 2θjet ), an 8-m class telescope with wideﬁeld imagers, such as the LSST and Subaru/HSC, is required. Oﬀ-axis orphan GRB afterglows up to ∼ 3θjet
can be discovered by performing time-domain surveys
with an 8-m class telescope. Because such optical timedomain surveys also detect numerous other optical transients (e.g. Urata et al. 2012), we presented expectedradio afterglow light curves for the conﬁrmation and determination of burst parameters. Radio light curves canbe monitored using ALMA and JVLA with reasonable
exposure (Figure 2 bottom).
1

Flux density [mJy]

Subaru Rc
Others R
Subaru Ic x 1.3
Others I x 1.3
Subaru z x 2.5
VLA 8.5GHz

0.1

Optical (r)On-axisiPTF
0.01

Θobs = 2Θjet

Pan-STARRS
Θobs = 3Θjet

10−3

HSC（3σ）

10−4

Θobs = 4Θjet
10−5

Θobs = 2ΘjetRadio (8.5GHz)

Θobs = 3ΘjetΘobs = 4Θjet

On-axisFlux density [mJy]

Flux density [mJy]

1

0.1

0.01

10−3

0.1

VLA（3σ）

1

10

100

Time [Days]
Fig. 2. Expected oﬀ-axis GRB afterglow with observing angles of2θjet (red), 3θjet (green), and 4θjet (blue) at typical redshift ofz = 0.5.

This work is partly supported by the Ministry of Science and Technology of Taiwan grants MOST103-2112M-008-021, 104-2112-M-008-011, and 105-2112-M-008013-MY3 (YU).
References
Urata, Y., et al. 2015, ApJ, 806, 222
Urata, Y., et al. 2012, ApJL, 760, L11

277

Gamma-ray burst observation for eight and half year by Akeno MITSuME
telescope
Taichi Fujiwara,1 Yoshihiko Saito,1 Ryosuke Itoh,1
Yutaro Tachibana,1 Yuki Ono,1 Kotaro Morita,1 Kouta Saisho1 and Nobuyuki Kawai1
1

Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro, Tokyo, Japan
E-mail(TF): fujiwara@hp.phys.titech.ac.jp
Abstract

We report the observations of early Gamma-Ray Burst (GRB) afterglows carried out with Akeno 50 cm
Telescope from 2008 to 2016. In this period, fourteen GRB afterglows were detected within 1000 seconds
after the detections of the prompt emission. We show the optical light curves of the GRB afterglows to
discuss timing properties of the external shock. Then, we compare the light curves with the forward/reverse
shock models. We discuss the number ratios of the diﬀerent light-curve types based on the simulation by
Gao et al. (2015). Finally, the ratio of magnetic equipartition parameters in the reverse and forward shocks
are evaluated as 10–100.
Key words: gamma-ray burst: general

1.

Introduction

The gamma-ray burst (GRB) afterglows are usually described well with forward shock emission. As the counterparts of the afterglows, reverse shock emission was
also predicted. However, the typical frequency of the reverse shock emission is lower than that of the forward
shock (Kobayashi et al. 2003).
Depending on the spectral peak frequency and the intensity ratio of the forward shock emission to that of thereverse shock, the early optical afterglows are classiﬁed
into four types (Gao et al. 2015):
Type 1: Rebrightening. The forward shock peak is beneath the reverse shock component. A sign of rebrightening emerges due to forward shock emission later.
Type 2: Flattening. The forward shock peak is beneath the reverse shock component. The forward shock
emission only shows its decaying part in the late stage.
Type 3: Forward shock-dominated emission without νm crossing. No reverse shock component is visible.
The observed optical peak is the deceleration peak.
Type 4: Forward shock-dominated emission withνm crossing. No reverse shock component is visible.
The observed optical peak is the crossing peak of νm .
It is necessary to observe GRB afterglows immediately
to analyze their morphology. Akeno 50 cm Telescope automatically observes GRBs detected by Swift and MAXI
to obtain the light curves of early afterglows.

Fig. 1. Optical light curves of GRB 111228A.

2. Observation and Results
From 2008 January to 2016 June, we observed 14 GRB
afterglows with Akeno Telescope within 1000 seconds after the detections of the prompt emission. Fig. 1 andFig. 2 show the examples of the light curves obtainedwith Akeno Telescope.
GRB 111228A showed a steep decay for < 200 seconds followed by ﬂattening. This light curve cannot be
ﬁt with a simple smooth broken power-law function but
with a combination of a single power-law component and278

Fig. 2. Optical light curve of GRB 100219A.

Fig. 4. Fitting result of the GRB 100219A data.

Fig. 3. Fitting result of the GRB 111228A data.

3. Discussion
Some GRBs cannot be classiﬁed, because the start oftheir observations were not early enough to detect theirrising phases if they exist.
We discuss number ratios of the diﬀerent light-curve
types by the simulation done by Gao et al. (2015). First,
if the reverse-to-forward shock ratio (RB ) of magnetic
equipartition parameters is generated through a Gaussian distribution with mean values (RB ) of 10, the cases
where the reverse shock dominates (Type 1 and 2) become less than 2%. We found 2 such bursts out of 14.
Therefore, RB needs to be larger than 10. If we assume
RB = 100, then the “reverse-shock dominant” cases become 20–30%, consistent with our results.
Thus, we ﬁnd that the reverse-to-forward shock ratioof magnetic equipartition parameters is likely 10 < RB ≤100.

a smooth broken power-law component. The model light
curve obtained in this ﬁtting is shown in Fig. 3. Thus,
we classify this afterglow as Type 1.
We detect an initial rise in the afterglow of GRB
100219A. We ﬁtted the light curve with a broken powerlaw function. The power-law index of the decaying slope
is smaller than 1.5, so we consider that a forward shock is
dominant in this GRB. We classify this GRB as a Type
3/4 candidate. The power-law index of the rising slope
is larger than 0.6, suggesting that the optical peak is thedeceleration peak. We therefore classify it into Type 3.
The result of this ﬁtting is shown in Fig. 4.
Applying a similar analysis to the early optical light
curves of the 14 GRBs, we eventually ﬁnd one burst ofType 1, another burst of Type 2, four bursts of Type 3
and eight bursts that cannot be uniquely classiﬁed.

References
Cano, Z. et al. 2014, A&A, 568, A19
Molinari, E. et al. 2007, A&A, 469, L13-16
Sari, R. 1997, ApJ, 489, L37
Costa, E. et al. 1997, Nature, 387, 783-785
Tanvir, R. et al. 2009, Nature, 461, 1254-1257
Zhang, B. et al. 2005, ApJ, 642, 354-370
Harrison, F. et al. 1999, ApJ, 523, L121
Tagliaferri, G. et al. 2005, A&A, 443, L1-5
Sari, R. et al. 1998, ApJ, 497, L17-20
Sari, R. et al. 2000, ApJ, 548, 787-799
Blanton, M. et al. 2007, AJ, 133, 734
Sesar, B. et al. 2006, AJ, 131, 2801-2825
Gao, H. et al. 2015, ApJ, 810, 160
Kobayashi, S. et al. 2003, ApJ, 582, L75-78
Nakar, E. et al. 2003, NewAstronomy, 8, 495-505

279

The ﬁrst Suzaku/WAM Gamma-ray Burst catalog
Norisuke Ohmori,1 Kazutaka Yamaoka,2,3 Masanori Ohno,4 Satoshi Sugita,5 Hiraku Yoshida,1Makoto Yamauchi,1 Kevin Hurley,6 Yuji Urata,7 and Makoto Tashiro8
1

Department of Applied Physics, University of Miyazaki, 1-1 Gakuen kibanadai-nishi,
Miyazaki-shi, Miyazaki 889-2192, Japan
2
Institute for Space-Earth Environmental Research (ISEE), Nagoya University, Furo-cho,
Chikuka-ku, Nagoya, Aichi 464-8601, Japan
3
Division of Particle and Astrophysical Science, Graduate School of Science, Nagoya University, Furo-cho,
Chikuka-ku, Nagoya, Aichi 464-8602, Japan
4
Department of Physics, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan
5
Department of Physics, Tokyo Institute for Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8550, Japan
6
Space Sciences Laboratory, University of California, 7 Gauss Way, Berkeley, CA 94720-7450, USA
7
Institute of Astronomy, National Central University, Chung-Li 32054, Taiwan
8Department of Physics, Saitama University, 255 Shimo-Okubo, Sakura-ku, Saitama-shi, Saitama 338-8570, Japan
E-mail(NO): ohmori@astro.miyazaki-u.ac.jp
Abstract
Gamma-ray bursts (GRBs) are one of the most powerful explosion phenomena. Many GRB catalogs
have been published to study the origin of the prompt emission so far, but the physical mechanisms remainunresolved. In order to investigate systematic properties of GRBs, we are constructing the GRB catalog for
GRBs observed by the wide-band all-sky monitor (WAM) onboard the Suzaku satellite. The WAM energy
range is 50 – 5000 keV, and its on-axis eﬀective area is large in comparison with other GRB instruments:
800 cm2 at 100 keV and 400 cm2 at 1 MeV. In the ﬁrst WAM GRB catalog, we will summarize resultsfrom 1464 GRBs detected between August 4, 2005 and December 29, 2010. Ohmori et al. (2016) alreadyreported the duration distributions with two distinct peaks at 0.38 sec and 19 sec, which conﬁrms thepresence of short/hard and long/soft GRBs. As a next step, we are performing spectral analysis with thethree models (power-law, power-law with an exponential cutoﬀ, and the Band function) using 412 GRBs
localized by Swift/BAT, Fermi/GBM and interplanetary network (IPN). In this paper, we report on thestatus of the catalog project, focusing on distributions of spectral parameters, ﬂuence, and 1-s peak ﬂux.
Key words: workshop: proceedings — gamma-ray burst: general — catalogs

1. Introduction
Gamma-ray bursts (GRBs) are one of the most powerful explosion phenomena. Since the Suzaku satellite waslaunched on 2005 July, the wide-band all sky monitor(WAM) succeeded in detecting many GRBs. The WAM
energy range is 50 – 5000 keV, and its on-axis eﬀective
area is large in comparison with other GRB instruments:
800 cm2 at 100 keV and 400 cm2 at 1 MeV (Yamaoka et
al. 2009). For now we will summarize results from 1464
GRBs detected between August 4, 2005 and December
29, 2010 in ﬁrst Suzaku/WAM GRB catalog. So far we
already reported the duration distributions. These results showed that the duration distributions with two
distinct peaks at 0.38 sec and 19 sec, which conﬁrms thepresence of short/hard and long/soft GRBs (Ohmori etal. 2016). Next we are performing spectral analysis us-

ing 412 GRBs localized by Swift/BAT, Fermi/GBM andinterplanetary network (IPN). In this paper, we reporton the status of the catalog project using 193 GRBs outof 412 GRBs, focusing on distributions of spectral parameters, ﬂuence, and 1-s peak ﬂux.
2. Analysis
In analysis, the spectra made using hxdmkbstspec andhxdmkwamspec in the HEASOFT version 6.19. The spectral ﬁtting tool used Xspec version 12.8. To ﬁt WAM
GRB spectra, we chose three models, which are powerlaw (PL), power-law with an exponential cutoﬀ (CPL),and the Band function (GRB) models(Band et al. 1993).
We ﬁtted total and 1-s peak spectra for each GRBs, andestimated energy ﬂuence and 1-s peak photon ﬂux.

280

Results

Table 2. Average and median list for each spectral parameter

GRB(short:long)
6 (0:6)

Figure 1 shows spectral parameter distributions of
CPL and GRB models for long and short GRB in diﬀerent detector. First, we compared WAM spectral parameter distributions between long and short GRBs. The
central value of α distributions for short GRB is shallower than that of long GRBs, which values are around-1 in long GRB and -0.6 in short GRB. This result showsthat short GRBs tend to have harder spectra than long
GRB. The average value of WAM β distributions in long
GRB is -2.83. More sample for GRB model are needed.
The WAM Epeak in long and short GRBs distributes
around 500 keV and 1 MeV, respectively.
Next, we also compared results with that of otherdetectors, which are Swift/BAT, CGRO/BATSE, andFermi/GBM. Table 2 shows average and median foreach spectral parameter distributions in diﬀerent detector(Lien et al. 2016; Goldstein et al. 2013; Gruber et al.
2014). These value shows that the α and β are similar
among detectors, but Epeak is so diﬀerent. These Epeak
distributions are strongly detector dependent.

2
1
500
200
100
50
20
10

2
1
0.5
0.2

BAT

2

WAM

20
10
5
2
10
5

BAT

1
0.5
100

BATSE

2
1
1000

BATSE

100

10
1
100

GBM

GBM

100

PLí3/2

0.1

1

10
100
Energy fluence (10í7 erg cmí2)

104

1000

Fig. 2. Integral distribution of WAM energy ﬂuence in 100 – 1000
keV. The lines show long GRB (red), short GRB (blue), and alldata (black). Green and light blue color lines show power-law witha slope of -3/2. These lines are drawn to guide the eye.
PLí3/2

BAT

100í1000 keVLong GRB
Short GRBAll

PLí3/2

BATSE

GBM
0.1

10

1

10

PLí3/2100í1000 keVLong GRB
Short GRB
All

10
1
1000

10

100

WAM

1230(910)1533(472)761(607)

WAM energy responce (∼20%) are not taken into accountin this analysis. This uncertainly will be estimated using
the cross-calibration results and published spectral parameters in the future.

100

2
1
10
5

WAM

-0.36(-0.59)-0.29(-0.54)-0.56(-0.58)

552(437)
140(102)
255(171)279(179)1

Number of GRBs

20
10
5

WAM
BAT
BATSE
GBM

100

CPL(short:long)54 (16:38)-0.80(-0.94)-1.07(-1.12)-1.02(-1.08)-0.94(-0.96)

Epeak [keV]
average(median)

10

PL(short:long)133 (25:108)

WAM
BAT
BATSE
GBM

β
average(median)long GRB-2.83(-2.74)-2.44(-2.43)-2.24(-2.22)
short GRB
-2.32(-2.44)-2.15(-2.34)

10

Model
Numberαaverage(median)

Number of GRBs

Table 1. Number list of the best ﬁt model

Detector

1

Table 1 shows number list of WAM spectral best ﬁtmodel. The WAM Epeak can be determinated by CPL
and GRB models for ∼31% out of 193 GRBs.

Number of GRBs

3.

1
10
1s peak photon flux (photons cmí2 sí1)

100

1
í2

0
alpha

2

í4

í3

í2
beta

í1

0

1

10 1001000 104
Epeak (keV)

Fig. 1. From left to right, the spectral parameter α, β, and Epeak
distributions for long GRB (red), short GRB (blue), and all data
(black). Also from top to bottom, WAM, BAT, BATSE, andGBM.

Figure 2 and 3 show integral distribution of ﬂuenceand 1-s peak ﬂux in 100 – 1000 keV. These value are
estimated from information of spectral best ﬁt model.
The energy ﬂuence distribution of long and short GRB
are spread from 4 × 10−7 to 2 × 10−4 erg cm−2 and 1 ×
10−7 to 7 × 10−5 erg cm−2 , respectively. Also, the 1-s
peak photon ﬂux distribution of long and short GRB are
spread from 0.4 to 40 photons cm−2 s−1 , and 0.4 to 15
photons cm−2 s−1 , respectively. The uncertainties of theFig. 3. Integral distribution of WAM 1-s peak photon ﬂux in 100 –1000 keV. The lines show long GRB (red), short GRB (blue), andall data (black). Green and light blue color lines show power-law
with a slope of -3/2. These lines are drawn to guide the eye.

References
Band, D. L., et al. 1993, ApJ, 413, 281
Goldstein, A., et al. 2013, ApJS, 208, 21
Gruber, D., et al. 2014, ApJS, 211, 12
Lien, A., et al. 2016, ApJ, 829, 7Ohmori, N., et al. 2016, PASJ, 68, S30
Yamaoka, K., et al. 2009, PASJ, 61, S35

281

SGRB rate from Yonetoku relation for SGRB and the detectability of GW
in O2 of aLIGO/aVirgo
Takashi Nakamura1
1

Kyoto University
E-mail(TN): nakamura.takashi.57a@st.kyoto-u.ac.jp
Abstract
Using 8 SGRB with Ep, z and Lp, Tsutsui et al. (2013) obtained Yonetoku relation for SGRB.
Applying this relation to bright BATSE SGRBs, the minimum event rate of SGRB is 6.3(+3.1−3.9 ) ×
10−10 events/Mpc3 /yr. Under the assumption of beaming factor of 100 and 4 times dimmer SGRBs,
if 10% of SGRB is NS-BH or 100% is NS-NS, there is a good chance to observe a few GW events in O2 of
aLIGO/aVirgo so that MAXI might see something at the same time.

282

GRB radio view from all angles
Alexander J. van der Horst1
1

George Washington University
E-mail(AV): ajvanderhorst@gwu.edu
Abstract
Gamma-ray bursts are among the most extreme objects in the Universe in terms of the energetics andthe physics needed to describe the observed phenomena. They are detected across the entire electromagnetic spectrum, from low-frequency radio waves to high-energy gamma rays. Our current understandingof gamma-ray bursts is based on multi-wavelength and multi-timescale observations. In this talk I willshow the role that radio observations have played, and will play in the future, in putting together a pictureof the physics behind the observed emission, the progenitors, and their environment. I will highlight theunique role that radio observations can play in detecting gamma-ray bursts for which the initial high-energy
emission is beamed away from us, so-called orphan afterglows, and the possibilities that several new andupgraded radio observatories oﬀer to obtain a better understanding of these enigmatic phenomena.

283

POLAR - The gamma-ray burst polarimetry experiment onboard China’s
TG-2 spacelab
Shuang-Nan Zhang1
1

Institute of High Energy Physics, China
E-mail(SZ): zhangsn@ihep.ac.cn
Abstract

On 2016-09-15, the gamma-ray burst polarimetry experiment, POLAR, onboard China’s Tiangong-2
spacelab has been launched. The main scientiﬁc goal of the experiment is to measure the polarizationof gamma-rays from gamma-ray bursts between 50–200 keV. POLAR will be able to detect about 50
gamma-ray bursts per year and can measure the polarization with precision of about 10% for the brightest
gamma-ray bursts. POLAR can also detect solar ﬂares and the Crab pulsar and measurement theirpolarization. In this talk I will describe the POLAR instrument and report some early results.

284

285

Hitomi

286

287

Hitomi observation of the Perseus cluster– gas motions and resonant scattering –Kosuke Sato,1 Maki Furukawa,1 Toru Sasaki,1 Kyoko Matsushita,1 Frits Paerels,2 Irina Zhuravleva,3Anna Ogorzalek,3 Norbert Werner,3 Steve Allen,3 Hiroya Yamaguchi,4 Greg Brown,4 Shinya Nakashima,5Yasushi Fukazawa,6 Masanori Ohno,6 Tetsu Kitayama,7 Brian McNamara,8and the Hitomi collaboration

1

Department of Physics, Tokyo University of Science, 2 Columbia University,
Department of Physics, Stanford University, 4 NASA/GSFC, 5 ISAS/JAXA,
6
Department of Physics, Hiroshima University, 7 Department of Physics, Toho University
8
Department of Physics and Astronomy, University of Waterloo
E-mail(KS): ksato@rs.tus.ac.jp
3

Abstract
Hitomi SXS (Hitomi; Takahashi et al. 2016, Proc. SPIE, 9905, 99050U, SXS; Kelley et al. 2016, Proc.
SPIE, 9905, 99050V) ﬁrst allow us to investigate ﬁne structures of emission lines from highly ionized ionswith its high energy resolution of ∼5 eV at 6 keV in orbit. The line diagnostics make it possible to measurestreaming and turbulent gas motions, which play key roles for understanding the basic physical propertiesof the intra-cluster medium (ICM) in galaxy clusters, such as the thermal conductivity and viscosity.
Particularly in the cluster core, resonance scattering should be taken into account when inferring physicalproperties from line intensities because the optical depth of the He-α resonant line is expected to be larger
than 1 (e.g., Zhuravleva et al., MNRAS, 2013).
Hitomi SXS carried out a series of 4 overlapping pointing observations of the Perseus cluster during
the Hitomi performance veriﬁcation phase with a total 300 ksec. Hitomi SXS ﬁnds an extremely quiescentatmosphere in the ICM of the Perseus cluster core where the gas has a line-of-sight velocity dispersion
below 200 km sec−1 (Hitomi collaboration, Nature, 2016). Also, the line-of-sight velocity map across the 60
kpc image of the cluster core shows a relative gradient of 150 ± 70 km sec−1 . These suggest that the cluster
mass measurement under the assumption of hydrostatic equilibrium would be little aﬀected by turbulent
pressure. On the other hand, the observed line ﬂux ratio of Fe XXV He-α resonant to forbidden lines isfound to be lower in the cluster core when compared to the outer region where the observed spectrum
shows a typical optically thin plasma feature, consistent with resonant scattering of the resonant line andalso in support of the low turbulent velocity. We also compare the observed line ﬂux ratios such as Fe He-αresonant, forbidden, He-β, and Ly-α with the expected ones from the Monte-Carlo simulation of resonant
scattering in the Perseus cluster core. The line ratios derived from the observation and simulations in thecluster core almost agree with each other within statistical and systematic errors.
Key words: galaxies: clusters: individual (Perseus cluster) – galaxies: clusters: intracluster medium– X-rays: galaxies: clusters288

Hitomi observation of the Perseus cluster– temperature and elemental abundances –Shinya Nakashima1 , Hiroya Yamaguchi2 , Kosuke Sato3
Kyoko Matsushita3 , Aurora Simionescu1 , Marshall Bautz4 ,and the Hitomi Collaboration
1

ISAS/JAXA, 2 NASA/GSFC, 3 Tokyo University of Science, 4 MIT/Kavli
E-mail(SN): shinya@astro.isas.jaxa.jp
Abstract

The X-ray emitting hot intracluster medium (ICM) dominates the baryonic mass in clusters of galaxies,
and hence the dynamics and physics of the ICM is crucial to understand the formation and evolution of thegalaxies and clusters. Several clusters have dense cores (the so-called “cool cores”) in which the radiative
cooling time of the ICM is shorter than the age of the universe. However, the observed temperature at
the core region is higher than expected, and moreover the star formation rate of the central galaxy is anorder of magnitude lower than that estimated from the cooling rate. Therefore a heating mechanism isrequired to suppress the cooling of the central dense ICM. Several heating sources are proposed but further
observations are necessary to constrain its origin. The metallicity of the ICM is also important because
it stores the metals ever ejected from cluster galaxies. Measurements of metal abundances in the ICM
provides a unique insight into the average star formation history in galaxies.
The Japan’s sixth X-ray astronomy mission, Hitomi (formerly ASTRO-H), was launched on 17th February 2016. The non-dispersive Soft X-ray Spectrometer (SXS) on board Hitomi achieves spectral resolutionof E/ΔE ∼ 1200 at 6 keV in orbits, which is about 30 times higher than that of X-ray CCDs. It provides
the ﬁrst opportunity to determine precise ionization sate, metal abundances, and turbulence velocity of
the ICM utilizing emission line diagnostics. Even though the operation of Hitomi was discontinued on
April 2016 because of the attitude anomaly, Hitomi observed the Perseus cluster for about 300 ks during
its commissioning phase.
We obtained an excellent spectrum of the ICM in the 2.0-20 keV band (limited energy range is due tothe closed gate valve in the commissioning phase). The He-like Fe Kα complex and the He-like Ni Kα lines
nearby the He-like Fe Kβ lines are successfully resolved. The weak emission lines from He-like Cr and Mn
are also clearly detected. The spectrum is almost represented by a single thin-thermal plasma model withthe temperature of ∼4 keV. However, the ionization temperatures of Si, S, Ar, Ca, and Fe derived fromdetailed line diagnostics deviates from a single temperature assumption, and therefore multi-temperatureor non-ionization equilibrium of the ICM is suggested. Indeed a two-temperature plasma model gives abetter ﬁt to the data. The precise plasma temperatures are under investigation because systematic eﬀects
due to the central AGN emission and the calibration uncertainty are not fully understood yet.
We also measured the abundance ratios of major elements to Fe assuming the single temperature plasma
model. All of them are ∼1.0 relative to the solar values, while previous CCD observations reported thesuper-solar abundances of Fe-peaked elements (Cr/Fe, Mn/Fe, and Ni/Fe). The cause of this discrepancy
is under discussion but we believe that our result utilizing the resolved Kα lines is reliable. The abundancesof Fe-peaked elements also shows the diversity of the progenitor of the type Ia supernova.
Key words: galaxies: clusters: individual (Perseus) — galaxies: clusters: intracluster medium —X-rays: galaxies: clusters

289

Hitomi Results of NGC 1275: The Origin of Fe-Kα Line
Hirofumi Noda1 , Yasushi Fukazawa2 , Koichi Hagino3 , Masanori Ohno2 , Masahiro Tsujimoto3 ,Takaaki Tanaka4 , Hiroyuki Uchida4 , Taiki Kawamuro4 , Francesco Tombesi5,6 , Scott Porter6 ,Caroline Kilbourne6 , Laura Brenneman7 , Christpher Reynolds5 , Richard Mushotzky5
and the Hitomi collaboration
1

Tohoku University (Japan), 2 Hiroshima University (Japan), 3 ISAS/JAXA (Japan),
Kyoto University (Japan), 5 University of Maryland (USA), 6 NASA/GSFC (USA),
7
Smithsonian Astrophysical Observatory (USA)E-mail(HN): hirofumi.noda@astr.tohoku.ac.jp

4

Abstract
A bright Fanaroﬀ-Riley I type radio galaxy NGC 1275 dominates the center of the Perseus cluster ofgalaxies at a redshift of z ∼ 0.01756 (Strauss et al. 1992, ApJS, 83, 29), and its radio lobes formed by thepast jet activity were found to cause kpc-scale cavities in the intracluster medium (ICM) by Chandra and
XMM-Newton observations (e.g., Fabian et al. 2006, MNRAS, 451, 3061). Hence, NGC 1275 is regarded
as a promising candidate to study important topics such as the cooling ﬂow (e.g., Fabian 1994, ARA&A,
32, 277) and interactions among the Active Galactic Nucleus (AGN), ICM, and galaxies (e.g., McNamara& Nulsen 2007 ARA&A, 45, 117).
During 2016 February–March, the Hitomi satellite (Takahashi et al. 2016, Proc. SPIE, 9905, 99050U)
observed the Perseus cluster including NGC 1275 with an X-ray microcalorimeter called the Soft Xray Spectrometer (SXS; Kelley et al. 2016, Proc. SPIE, 9905, 99050V) for the ﬁrst time (The Hitomi
collaboration, 2016, Nature, 535, 117). In the present study, we analyz the Hitomi datasets and archival
datasets of other X-ray satellites, mainly aiming to reveal X-ray generation mechanism and the structureof the AGN region in NGC 1275. One of the most important mysteries is the origin of the Fe-Kα emission
line at 6.4 keV (Churazov et al. 2003, ApJ, 590, 225; Yamazaki et al. 2013, PASJ, 65, 30). Possible sitesof origin include an outer region of an accretion disk, a Broad Line Region (BLR), a dusty torus, and arotating molecular disk toward outside within the AGN system. In addition, NGC 1275 was reported tohave a few tens kpc-scale molecular clouds and Hα ﬁlaments outside the AGN (e.g., Salomé et al. 2006,
A&A, 454, 437), and hence, they are also candidates of the Fe-Kα emitter.
By utilizing the unprecedented energy resolution of ∼ 5 eV at 6 keV achieved by the SXS, we signiﬁcantly detect the Fe-Kα line with an equivalent width of ∼ 10 eV against a total X-ray continuum
from the AGN and ICM (∼ 25 eV against only the AGN continuum). We successfully limit its velocity
width to ∼ 500–1400 km s−1 (FWHM) for the ﬁrst time ever. Because the derived Fe-Kα velocity widthis signiﬁcantly narrower than that of broad Hα of ∼ 2750 km s−1 (Ho et al. 1997, ApJS, 112, 391), we
can exclude a large contribution of Fe-Kα emission from an the accretion disk and the BLR. Furthermore,
we perform Monte Carlo simulations which calculate the Fe-Kα intensity from the molecular clouds byutilizing the MONACO framework (Odaka et al. 2011, ApJ, 740, 103), and ﬁnd that their contribution is
also too small. Therefore, we suggest that the origin of the Fe-Kα line from NGC 1275 is likely from thedusty torus to the rotating molecular disk which extends to a few hundred pc scale (e.g., Scharwächter et
al. 2013, MNRAS, 429, 2315). Further limitation of the Fe-Kα source in NGC 1275 and comparison withother type AGNs will be discussed in a forthcoming paper.
Key words: galaxies: active – galaxies: nuclei – galaxies individual (NGC 1275)

290

291

Hitomi Observations of Supernova Remnants N132D and G21.5-0.9
Satoru Katsuda,1 Hiroya Yamaguchi,2,3 Makoto Sawada,4 Eric D. Miller,5 Kumiko Nobukawa,6Masayoshi Nobukawa,7 Takaaki Tanaka,8 Hiroshi Nakajima,9 Aya Bamba,10 John P. Hughes,11Masahiro Tsujimoto,12 Toshiki Sato,12,13 Kosuke Sato,14 Lorella Angelini,2 Benson Guest,15
Hiroyuki Uchida,8 Hiromitsu Takahashi,16 Yukikatsu Terada,17 Shu Koyama,12 Kenya Oshimizu,10Caroline A. Kilbourne,2 Katsuji Koyama,8 Samar Saﬁ-Harb,15 Brian J. Williams,2and the Hitomi Collaboration1

Chuo University (Japan), 2 NASA GSFC (USA), 3 University of Maryland (USA),
Aoyama Gakuin University (Japan), 5 MIT (USA), 6 Nara Womens University (Japan),
7
Nara University of Education (Japan), 8 Kyoto University (Japan), 9 Osaka University (Japan),
10
University of Tokyo (Japan), 11 Rutgers University (USA), 12 JAXA ISAS (Japan),
13
Tokyo Metropolitan University (Japan), 14 Tokyo University of Science (Japan),
15
University of Manitoba (Canada), 16 Hiroshima University (Japan), 17 Saitama University (Japan)
E-mail(SK): katsuda@phys.chuo-u.ac.jp
4

Abstract
We present preliminary results from Hitomi observations of supernova remnants (SNRs) N132D and
G21.5-0.9 which were selected for the performance veriﬁcation as well as initial calibration of the Hitomi
satellite.
N132D is the X-ray brightest SNR in the Large Magellanic Cloud (LMC). It is a member of a rareclass of oxygen-rich young SNRs, characterized by the presence of fast-moving oxygen-rich optical knots
thought to result from He-burning nucleosynthesis in the cores of massive stars. In the early phase of the
Hitomi observation of N132D, the satellite attitude control was unfortunately lost due to problems in thestar tracker system (STT). Therefore, the exposure time of the soft-X-ray spectrometer (SXS) observation
was unexpectedly short (less than an hour). Nevertheless, the SXS successfully detected some K-shell lines
from highly-ionized sulfur and iron, thanks to its excellent spectral resolution and gain accuracy. On theother hand, the soft X-ray imager (SXI) was able to catch N132D even after the STT anomaly, thanks
to its large ﬁeld of view. Therefore, we obtained a statistically-rich SXI spectrum. We also investigated
X-ray spectra taken by the Suzaku XIS and the NuSTAR FPM, and found clear signatures of Fe Lyα
and Fe Heβ (Bamba et al. 2017 in preparation). Based on our modeling of the combined SXI+XIS+FPM
spectra, we revealed that the Fe-K emitting plasma is recombining, and that the SXS-detected Fe-K lines
complex, namely the strongest Fe-K feature in the X-ray spectrum, is likely to be Fe Heα lines. Using the
SXS spectrum, we found these Fe Heα lines to be redshifted by ∼30 eV or ∼1500 km s−1 . This velocity ismuch larger than the LMC’s systemic velocity along the line of sight (278 km s−1 : Mateo 1998, ARA&A,
36, 43), which can be naturally explained by an asymmetric supernova explosion. In addition, we wereable to measure a line of sight velocity of the S Heα lines to be ∼5 eV or ∼600 km s−1 , roughly consistent
with the LMC’s systemic velocity. It should be noted, however, that the SXS energy-scale uncertainty inthe low-energy band is larger than that around Fe Heα, and may be equivalent to the shift measured for
S Heα. Cautious investigations of the SXS energy scale are now in progress to better constrain the S Heαlines’ centroid.
G21.5-0.9 is a plerionic SNR with a bright Crab-like pulsar wind nebula (PWN) in our Galaxy, having
a ﬁlled synchrotron nebula centered on the ∼62 ms pulsar, PSR J1833-1034. Hitomi observed G21.5-0.9
with the SXS, SXI, and the hard X-ray imager (HXI), providing a high-statistics wide-band (∼1–60 keV)
spectrum. The wide-band spectrum can be statistically better ﬁtted with a broken power-law model thana simple power-law model, consistent with the previous NuSTAR result. This indicates either a break in
the injected electron energy spectrum or energy losses due to particle transport in the PWN (Nynka et al.
2014, ApJ, 789, 72). The pulsation period and the period derivative of PSR J1833-1034 were measured
in radio and GeV gamma-rays to be P = 61.86 ms and Ṗ = 2.0 × 10−13 s s−1 (Camilo et al. 2006, ApJ,
637, 456; Abdo et al. 2013, ApJS, 208, 17), but have not yet been detected in the X-ray band. So far, ourtiming analysis using both the HXI and the SXS does not show any signiﬁcant pulsations partly due tothe much brighter PWN. However, given that the pulsed fraction (if any) would be fairly small in X-rays,
we need further careful analyses of systematic uncertainties to ﬁnd possible pulsations or give a stringent
constraint on a pulsed fraction. In addition to these studies, a survey of emission and absorption lines in
the SXS spectrum is ongoing, which will be presented in a forthcoming paper.
Key words: ISM: individual objects (N132D, G21.5-0.9) — ISM: supernova remnants — X-rays: ISM

292

Atomic data and spectral modeling constraints from high-resolution X-ray
spectroscopic observations of the Perseus cluster with Hitomi
M. Sawada1 , J. Kaastra2 , 3, L. Gu2 , H. Akamatsu2 , J. de Plaa2 , C. de Vries2 , R. Smith4 , A. Foster4 ,H. Odaka5 , M. Leutenegger6 , M. Markevitch6 C. Kilbourne6 , H. Yamaguchi6 , G. Brown7 , F. Paerels8 ,N. Ota9 , T. Tamura10 , S. Nakashima10 , S. Ueda10 , K. Matsushita11 , K. Sato11 ,R. Mushotzky12 , T. Ohashi13 , and Hitomi Collaboration
1

Aoyama Gakuin University, 2 SRON, 3 Leiden University, 4 CfA, 5 KIPAC, 6 NASA/GSFC, 7 LLNL,
Columbia University, 9 Nara Women s University, 10 ISAS/JAXA, 11 Tokyo University of Science,
12
University of Maryland, 13 Tokyo Metropolitan University
E-mail(MS): sawada@phys.aoyama.ac.jp

8

Abstract
High-resolution X-ray spectroscopy is a key to understand nature of cosmic plasmas, as it provides direct
measurements of temperature and velocity structures, chemical abundances, and charge-state distributions.
It also opens a window for new spectral features. Plasma codes are indispensable to carry out highresolution spectroscopy as any spectral analysis uses a plasma code to derive physical parameters. The ﬁrst
high-resolution X-ray spectrum of the Perseus cluster has been obtained with the Soft X-ray Spectrometeron the Hitomi X-ray satellite. It covers major spectral features in 2–9 keV, including K-shell emission lines
of abundant elements Si through Ni. The Hitomi spectrum can be used for the ﬁrst practical benchmarkof plasma codes. One of our purposes is to evaluate systematic uncertainties in calculations of spectra andtheir impacts on deriving physical parameters. Another purpose is to investigate possibilities to detect
new spectral features, for instance charge-exchange (CX) emission or characteristic emission line ratios
due to interactions with non-thermal electrons.
Starting from the same dataset as presented by Hitomi Collaboration et al. (2016a, b), empiricalcorrections for energy scale, eﬀective area, and low-energy redistribution matrix were ﬁrst applied. The
baseline model was deﬁned based upon the corrected data with the SPEX package version 3.03.03. It
is mainly a hot gas component of 4-keV plasma at collisional ionization equilibrium (CIE) aﬀected bypossible resonance scattering, interstellar absorption, and spectral broadening due both to thermal andturbulent motions. The other component is emission from the central active galactic nucleus. It consists ofa power-law continuum and Gaussians for the neutral Fe-K lines, whose parameters have been kept frozen
throughout the analysis.
Systematic uncertainties of the spectral parameters were estimated as deviations from the best-ﬁt values
with the baseline model. Among the systematic eﬀects considered in our analysis, the choices of plasma
codes and atomic data behind codes showed major impacts. In the comparison of the codes, it is foundthat the relative diﬀerence between SPEX and APEC is usually within 10%. The origins depend on thelines, either diﬀerent line wavelengths or diﬀerent emission rates. An emission rate is determined by a
combination of charge-state distribution and transition rate. On transition rates, we take the He-like Fe-K
multiplet lines as a test case to assess errors. For the resonance and inter-combination lines, predominantly
formed by electron impact excitation from the ground state and radiative decay back there, the rates by
the major databases (SPEX, APEC, and CHIANTI) and the theoretical calculation by the FAC agree wellwith each other within 10%. On the other hand, for the forbidden line, the relative diﬀerence is found tobe as large as 40%, likely because of its complexity having multiple forming channels.
New spectral features are also under investigation. One is the CX emission from the colliding surface
between hot and cool gas components, characterized by high-shell transition lines. With the baselinemodel, we found two line-like residuals near at 3.4 and 8.6 keV. These energies are consistent with high-shelltransition lines of S XVI and Fe XXV, respectively, although the signiﬁcances are below 3σ. The presence of
non-thermal electrons can also modify line emission for instance by changing inner-shell transition rates.
By using a capability of the SPEX CIE model to handle this, we obtained upper-limits of the relative
density to thermal population as a function of the mean energy, with a quasi-thermal assumption where
the non-thermal distribution is approximated by a Maxwellian.
All the results are currently being reviewed and updated with the ﬁnal version of the pipeline products,
software tools, and calibration database, and will be published in a full separate paper.
Key words: : atomic data — atomic processes — clusters: individual (Perseus) — galaxies: intergalactic medium — X-rays: galaxies: clusters

293

Hitomi Observation of the Highly Obscured High-Mass X-ray Binary
IGR J16318–4848
Hiroshi Nakajima,1 Kiyoshi Hayashida,1 Tim Kallman,2Caroline Kilbourne2 and the Hitomi Collaboration
1

Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka, 560-0043, Japan
2
NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
E-mail(HN): nakajima@ess.sci.osaka-u.ac.jp
Abstract

IGR J16318–4848 is a high-mass X-ray binary (HMXB) system with an extremely strong absorption.
It was discovered by INTEGRAL (Courvoisier et al. 2003, IAUC, 8063) and the X-ray spectrum wasfound to be dominated with strong ﬂuorescence lines of iron and nickel as well as the continuum emission(Matt & Guainazzi 2003, MNRAS, 341, L13). Long term light curves by RXTE-ASM, Swif t-BAT and
INTEGRAL-ISGRI suggest a possible orbital period of 80 d (Jain et al. 2009, RAA, 12, 1303). The
ﬂuorescence lines and the continuum vary on time scales of thousands of seconds, corresponding to an
emitting region of about 1013 cm (Walter et al. 2003, A&A., 411, L427). The optical/near-infrared(NIR) counterpart exhibits less absorption than that measured in the X-ray band, which implies thatthe absorbing material is concentrated around the compact object. The NIR spectrum also suggests thecounter part is a supergiant B[e] star (Filliatre & Chaty 2004, ApJ, 616, 469) and hence the source isclassiﬁed as a HMXB system. However, the nature of the compact source (neutron star or blackhole) is
not yet conﬁrmed.
Hitomi observation was performed just after the startup of the Soft X-ray Spectrometer (SXS), the Soft
X-ray Imager (SXI), and one of the two Hard X-ray Imager (HXI). Because alignment matrices of the star
trackers were not yet optimized, the source was outside of the ﬁeld-of-view (FoV) of SXS for most of the
exposure. On the other hand, larger FoVs of HXI and SXI enabled us to obtain a broad band spectrum
and light curve. The eﬀective exposure time of these two instruments are 42 ks and 143 ks for HXI andSXI respectively, although we use the data only when both of the two instruments were available in thecurrent analyses.
SXT+SXI and HXT+HXI spectra show negligible contribution of the Compton shoulder as in theprevious observations (Ibarra et al. 2007, A&A, 465, 501). The spectra are ﬁtted with an absorbed cut-oﬀ
power-law and three gaussians to account for the ﬂuorescence lines. Comparing the spectral parameterswith those obtained from Suzaku observation in 2006 (Barragan et al. 2009, A&A, 508, 1275), the ﬂux
of continuum and line components signiﬁcantly decreased in this ten years while the equivalent widths
increased. Unabsorbed luminosity in 2 to 10 keV is 5.8×1035 ergs/s, which is far below the Eddington
limit of 1.8×1038 ergs/s for a neutron star of 1.4 M and hence permits moderate accretion. The ratio of
the continuum ﬂux to the ﬂuorescence line ﬂux is positively correlated with the continuum ﬂux. In otherwords, the continuum component varies more dynamically than the line component, or at least a part ofthe line emission is not aﬀected by the continuum emission.
In spite of a severely limited number of source events, we successfully detected Fe Kα1 and Fe Kα2 lines
from SXS spectrum thanks to its outstanding energy resolution. We put a constraint on the Kα1 centroid+1.7
to be 6404.8 −1.3
eV, which is tightest among the past observations. Combining the best-ﬁt parameters
obtained with SXI and HXI, neutral or slightly ionizing state of Fe is plausible, although the gain uncertainty of each instrument should be referred. Considering the line-of-sight velocity of the line-emitting
matter, NIR spectrum is consistent with no systemic velocity of the sgB[e] star (cΔλ/λ=110±130 km/s).
Assuming that the line-emitting matter follows the compact object that rotates around a companion witha mass of 30 M , a distance of 1013 cm, and an inclination angle of 45◦ , the maximum line-of-sight velocity
will be 450 km/s (9 eV), which does not change our estimation.
Coherent pulsation search is performed both for SXI and HXI light curves from 1 s to one tenth of the
exposure time of each instrument. We ﬁnd a hint of periodicity at 46.9 s in the HXI light curve for theenergy band from 5.5 to 80 keV.
Key words: Stars: individual: IGR J16318–4848 — binaries: spectroscopic — X-rays: binaries

294

295

The Hard X-ray Imager onboard Hitomi (ASTRO-H)
Kazuhiro Nakazawa,1 Goro Sato,2 Motohide Kokubun,2 Teruaki Enoto,3 Yasushi Fukazawa,4 Kouichi Hagino,2Atsushi Harayama,2 Katsuhiro Hayashi,5 Jun Kataoka,6 Junichiro Katsuta,4 Philippe Laurent,7 Francois Lebrun,7Olivier Limousin,7 Kazuo Makishima,8 Tsunefumi Mizuno,9 Kunishiro Mori,2 Takeshi Nakamori,10 Toshio Nakano,11
Hirofumi Noda,12 Hirokazu Odaka,13 Masanori Ohno,4 Masayuki Ohta,2 Shinya Saito,14 Rie Sato,2 Hiroyasu Tajima,15
Hiromitsu Takahashi,4 Tadayuki Takahashi,2 Shin’ichiro Takeda,16 Yukikatsu Terada,17 Hideki Uchiyama,18
Yasunobu Uchiyama,14 Shin Watanabe,2 Kazutaka Yamaoka,15 Yoichi Yatsu,19 Takayuki Yuasa,8 and the HXI team1

Department of Physics, University of Tokyo, Tokyo 113-0033, Japan
2
ISAS/JAXA, Sagamihara, Kanagawa 252-5210, Japan
3
Faculty of Science, Kyoto University, Kyoto 606-8502, Japan
4
Department of Physical Science, Hiroshima University, Hiroshima 739-8526, Japan
5
Division of Particle Physics and Astronomy, Nagoya University, Aichi 464-8602, Japan
6
Research Institute for Science and Engineering, Waseda University, Tokyo 169-8555, Japan
7
APC/ IRFU / Service d’Astrophysique, CEA Saclay, 91191 Gif-sur-Yvette Cedex, France
8
MAXI team, Global Research Cluster, RIKEN, Saitama 351-0198, Japan
9
Hiroshima Astrophysical Science Center, Hiroshima University, Hiroshima 739-8526, Japan10
Department of Physics, Yamagata University, Yamagata 990-8560, Japan
11
Nishina Center, RIKEN, Saitama 351-0198, Japan
12
The Frontier Research Institute for Interdisciplinary Sciences, Tohoku University, Sendai 980-8578, Japan
13
SLAC, National Accelerator Laboratory, Stanford University, California 94025-7015, USA
14
Department of Physics, Rikkyo University, Tokyo, 171-8501, Japan
15
Institute for Space-Earth Environmental Research, Nagoya University, Aichi 464-8601, Japan
16
Advanced Medical Instrumentation Unit, OIST, Okinawa, 904-0495 Japan17
Department of Physics, Saitama University, Saitama 338-8570, Japan
18
Science Education, Faculty of Education, Shizuoka University, Shizuoka, 422-8529, Japan
19
Department of Physics, Tokyo Institute of Technology, Tokyo 152-8551, Japan
E-mail(FS): nakazawa@juno.phys.s.u-tokyo.ac.jp
Abstract
The hard X-ray imager (HXI) onboard Hitomi (ASTRO-H) mission was coupled with the hard X-ray
telescope (HXT) to provide ﬁne imaging spectroscopy at 5 to 80 keV. The system is made of 2 sets of
HXI+HXT, covering a ﬁeld-of-view of 9×9 armin2 . Hitomi was launched in 17 February 2016 and the HXIstarted observation on 14 March, and was operational till the incident lead to Hitomi loss on 26 March. All
detector channels worked well and showed performance consistent with those seen on ground. Calibration
observations conﬁrmed the 5–80 keV energy coverage, with eﬀective area meeting the requirement. The
background level was shown to be as low as 1–3×10−4 counts s−1 keV−1 cm−2 (here cm2 is in detector
geometrical area), also meeting the requirement. These properties showed that the HXI was one of themost sensitive instruments in 10-80 keV.
Key words: Hitomi, Hard X-ray Imager, DSSD, CdTe-DSD, BGO, in-orbit performance

1. INTRODUCTION
Hitomi X-ray observatory (e.g. Takahashi et al., 2016)
launched in 17 February 2016 had a hard X-ray imaging
spectroscopy system made of two hard X-ray imagers(HXIs, Nakazawa et al., 2016) coupled with two hard Xray telescopes (HXTs, Awaki et al. 2016). With 12 m
focal length, they provide ﬁne (1 .7 half-power diame-ter) imaging spectroscopy at 5 to 80 keV. They are designed to provide crucial information about hottest thermal emission, non-thermal emission, as well as reﬂection
components of high-energy astrophysical objects.
Speciﬁcation of the HXI+HXT system is summarizedon Table 1 and Figure 1. Two HXIs are mounted on
the HXI plate, connected to the satellite baseplate via

296

100

NuSTAR FPMA
HXI-1 after layer selection

4 arcmin radius integrated
effective area
5

10

20

50

100

Energy (keV)
2

80 arcmin scaled background
NuSTAR FPM

HXI-1 after layer selection
10 4

the Extensible Optical Bench (EOB). The EOB deploysby 6 m after launch. On top of the satellite, the Fixed
Optical Bench is mounted and the two HXTs are locatedon its top-plate. With this conﬁguration, the 12 m focal
length of the HXTs is achieved.
The HXI imagers are made of 4 layers of Si and a
CdTe semiconductor double-sided strip detectors, with32 × 32 mm size, stacked to enhance detection eﬃciency
as well as to enable photon interaction-depth sensing,valid for background reduction (G. Sato et al. 2016).
Active shield made of 9 BGO scintillators surrounds theimager to provide low background environment. Afterthe EOB deployment on 28 February, the HXI was gradually turned on. Two imagers successfully started observation on 14 March, and was operational till 26 March.

HXI-1 total
Effective Area (cm2)

Fig. 1. (left) The hard X-ray imaging spectroscopy system of Hitomi. The 2 HXI-Ss are located at the HXI plate, connected tothe satellite main body via the EOB. (right) Photo of the HXIplate equipped with 2 HXI-Ss covered with thermal blanket using
Ag-Teﬂon layer at the surface. Also seen in front is the radiator.

10

HXI plate

0.1

2x HXIs

0.01

Extensible Optical
Bench (+6m)

10 3

12 m
focal length

sitive to ∼ 100 keV electrons in orbit, after proper data
selection, the background level was shown to be as low
as 1–3×10−4 counts s−1 keV−1 cm−2 (here cm2 is in
detector geometrical area). As shown in Figure 2, the
HXI background is lower, and the eﬀective area is higher
than those of NuSTAR. Combined with ∼ 2 times worse
angular resolution, we can see that the statistical sensitivity limit is similar for point source in the two missions.
For largely extended sources, the HXI+HXT’s statisticallimit is twice better. Because of the nearly perfect CXB
baﬄing, which could not be achieved in NuSTAR, the
HXI has more stable background and better potential
for dim extended sources, as well.

normalized counts s 1 keV 1

2x HXTs

5

10

20

50

Energy (keV)

Table 1. Speciﬁcation of the HXI+HXT system (2 systems added).

angular resolution
FoV
eﬀective area
energy band1 .7@30 keV HPD
9 .17 × 9 .17 (32 × 32 mm)> 300 cm2 @ 30 keV total
5–80 keV

2. In-orbit performance
After the commissioning operation, we ﬁrst veriﬁed all
1280 ch each of the camera and 9+2 ch each of the APDs
were working well. Hitomi observed a pulsar wind nebula G21.5−0.9 as the ﬁrst light for the HXI. Then, it
observed an isolated neutron star RXJ1856−3754 forsoft X-ray instruments. Then, Hitomi observed Crabfor 4 orbits, i.e. ∼6 hours in duration. The Crab andG21.5−0.9 spectra veriﬁed the eﬀective area of the system (see Awaki et al. 2016).
The RXJ1856−3754 data provided us with the background data. Although the HXI turned out to be sen-Fig. 2. (left) Preliminary eﬀective area accumulated over a radius of
4 of the HXI1-S. Also shown in black in comparison is that of
the NuSTAR FPMA. (right) Preliminary background level of the
HXI1-S, after the layer selection accumulated over 80 arcmin2 area
(magenta), equivalent to 9.9 cm2 geometrical area. Also shown(in black) in comparison is that of the NuSTAR FPM, scaled fromthe NuSTAR proposal information (6.9 cm2 geometrical).

References
Takahashi, T. and Hitomi-collaboration, “The astro-h xray astronomy satellite”, Proc .SPIE 9905 (2016).
Awaki, H. et al., “Performance of astro-h hard x-ray
telescope (hxt)”, Proc .SPIE 9905 (2016).
Nakazawa, K. et al., “The Hard X-ray Imager onboard
ASTRO-H”, Proc .SPIE 9905 (2016).
Sato, G. et al., “The Si/CdTe semiconductor cameraof the ASTRO-H Hard X-ray Imager (HXI)”, NIM-A
831, p.235–241 (2016).

297

In-ﬂight performance of the Soft Gamma-ray Detector (SGD) onboard
Hitomi
Yasushi Fukazawa,1,2 Hiroyasu Tajima,3 Shin Watanabe,4,5 Roger Blandford,6 Teruaki Enoto,7Andrea Goldwurm,8 Kouichi Hagino,4 Katsuhiro Hayashi4,9 Yuto Ichinohe,10 Jun Kataoka,11Junichiro Katsuta,1 Takao Kitaguchi,1 Motohide Kokubun,4 Philippe Laurent,8 François Lebrun,8Olivier Limousin,8 Grzegorz M. Madejski,6 Kazuo Makishima,12 Tsunefumi Mizuno,2 Kunishiro Mori,4Takeshi Nakamori,13 Toshio Nakano,12 Kazuhiro Nakazawa,5 Hirofumi Noda,14 Hirokazu Odaka,6Masanori Ohno,1 Masayuki Ohta,4 Shinya Saito,15 Goro Sato,4 Rie Sato,4
Shin’ichiro Takeda,16 Hiromitsu Takahashi,1 Tadayuki Takahashi,4,5 Takaaki Tanaka,7 Yasuyuki Tanaka,2
Yukikatsu Terada,17 Hideki Uchiyama,18 Yasunobu Uchiyama,15 Kazutaka Yamaoka,2,9 Yoichi Yatsu,19Daisuke Yonetoku,20 Takayuki Yuasa,121

Department of Physical Science, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan
Hiroshima Astrophysical Science Center, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526,
Japan
3
Institute for Space-Earth Environmental Research, Nagoya University, Nagoya, Aichi 464-8601, Japan
4
Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, Sagamihara,
Kanagawa 252-5210, Japan
5
Department of Physics, University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan
6
KIPAC, Stanford University, Stanford, CA 94305, USA
7
Department of Physics, Kyoto University, Kitashirakawaoiwake, Sakyo Kyoto, Kyoto, 606-8502, Japan
8
IRFU / Service d’Astrophysique, CEA Saclay, 91191 Gif-sur-Yvette, Cedex France
9
Department of Physics, Nagoya University, Nagoya, Aichi 464-8601, Japan10
Department of Physics, Tokyo Metropolitan University, Hachioji, Tokyo 192-0397, Japan
11
Research Institute for Science and Engineering, Waseda University, Shinjuku-ku, Tokyo, 169-8050,
Japan
12
Cosmic Radiation Laboratory, Institute of Physical and Chemical Research (RIKEN) Wako, Saitama
351-0198, Japan
13
Department of Physics, Faculty of Science, Yamagata University, Kojirakawa, Yamagata 990-8560,
Japan
14
Frontier Research Institute for Interdisciplinary Sciences, Tohoku University, Sendai, Miyagi
980-8578, Japan
15
Department of Physics, Rikkyo University, Toshima-ku, Tokyo, 171-8501, Japan
16
Advanced Medical Instrumentation Unit, Okinawa Institute of Science and Technology Graduate
University, Onna, Okinawa 904-0495, Japan17
Department of Physics, Saitama University, Saitama, Saitama 338-8570, Japan
18
Science Education, Faculty of Education, Shizuoka University, Suruga-ku, Shizuoka, 422-8529,
Japan
19
Department of Physics, Tokyo Institute of Technology, Meguro-ku, Tokyo 152-8551, Japan20
Department of Physics, Kanazawa University, Kanazawa, Ishikawa 920-1192, Japan
E-mail(YF): fukazawa@astro.hiroshima-u.ac.jp

2

Abstract
The Soft Gamma-ray Detector (SGD) is one of observational instruments onboard the Hitomi (ASTROH), and will provide better sensitivity in 60–600 keV than the past observatories. The SGD utilizes similar
technologies to the Hard X-ray Imager (HXI) onboard the ASTRO-H. The SGD achieves low backgroundby constraining gamma-ray events within a narrow ﬁeld-of-view by Compton kinematics, in addition to the
BGO active shield. SGD was successfully turned on and observed the Crab nebula. All the instrumental
components, including 6 Compton cameras and 50 BGO active shields, were worked well without signiﬁcant
problems. Noise performance was <2 keV (FWHM) at 100 keV, as well as expected, and the pulse proﬂe
from the Crab netron star was obtained. Polarization studies of the Crab nebula is going on. In this paper,
we will report the in-ﬂight operation, calibration, and performance of SGD.
Key words: Gamma-rays Instruments — Gamma-rays: Soft Gamma-ray Compton Camera

298

103

511 keV peak
FWHM 6.1-7.1 keVCounts/bin(0.75 keV)Counts/bin(0.75 keV)

336 keV peak
FWHM 4.2-4.5 keV

400
350

511 keV peak

300

FWHM 6.6 keV
250
200

102

150
100
50

100

200

300

400

500

600
700
Energy[keV]

0
0

100

200

300

400

500

600
700
Energy[keV]

Fig. 1. SGD1 CC1 spectra in orbit. The left panel shows the spectraof the CdTe detectors. The sum spectrum of 48 side CdTe detectors and that of 32 bottom CdTe detectors are shown in blackand red, respectively. The right panel shows the spectrum ofCompton-reconstructed events. Gamma-ray emission lines from
activated materials can be seen.

value, implying that the eﬀects of the ﬁne collimator
misalignment and distortion are negligible.
In order to conﬁrm the time assignment, pulse analysesof the Crab pulsar were performed. In order to obtainlarge number of events, photo-absorption signals in the
top 4 Si layers are used in this analysis. The tool for
ASTRO-H time assignments is applied to the SGD Crabdata, and, barycentric correction is also applied. The
pulse period of 33.720462 ms was successfully detected,
as shown in Figure 2 left.
Pulse Profile
Counts/sec/keV

Compton reconstruction spectrum

Bottom CdTe (32 devices)

CountsSide CdTe(48 devices)

1200
1100

102

1000
900

103

800

A Compton camera of SGD consists of a stack of 32
layers of 0.6 mm thick Si pixel sensors followed by 8
layers of 0.75 mm thick CdTe pixel, surround by 2 layers of 0.75 mm thick CdTe pixel sensors on the foursides. Signal processing chain of the Compton camera
consists of 208 Front-End Cards (FECs) followed by four
ASIC Driver Boards (ADBs) and an ASIC Control Board(ACB). These components are packed in a 12×12×12
cm3 aluminum enclosure. Such a high-density stackingof Si and CdTe sensors is critical to obtain a high eﬃciency of Compton reconstruction of 10–15% around 100
keV. Detailed design of Compton camera of SGD is described in Watanabe et al. (2014).
The start-up operation of SGD-S began from March
15, 2016. From March 21, SGD1-S had been operated in
the nominal observation mode, and SGD2-S was put into
the nominal observation on March 24 2016, just before
the maneuver to the Crab nebula. The exposure of the
Crab observation was about 5 ks. The data before themaneuver to the Crab nebula is useful for the backgroundstudies of SGD.
In orbit, we conﬁrmed the performance by using
gamma-ray emission lines from activated materials. Figure 1 shows the single-hit spectra of CdTe sensors andthe spectrum of Compton reconstructed events in orbit.
Good energy resolutions consistent with the expectationfrom the on-ground performance were also veriﬁed as
there was no degradation from the performance obtainedin the component level tests.
In order to conﬁrm the alignment with the optical axisof the satellite, we compared the observed rate for the
Crab nebula with the Monte Carlo simulation result assuming the nominal Crab ﬂux. The count rates of 35–
60 keV, where the eﬀects of threshold dispersions are
insigniﬁcant, were obtained within less than 5 percents,
which means that the misalignments among six ﬁne collimators was less than 1.5 arcmin. The count rate predicted by simulation was consistent with the observed

700
600
104

500
0

0.2

0.4

0.6

0.8

1

1.2

1.4

1.6

1.8
2
Phase

50

60 70 80

100

200

300
Energy[keV]

Fig. 2. (left) Obtained pulse proﬁle of top 4 layers Si detector events.
All events of six Compton cameras are summed. The pulse period is 33.720462 ms. (right) Compton event reconstructed spectrum during the Crab observation. The background is subtracted.
The events from four Compton cameras (SGD1CC1,CC2,CC3 and
SGD2CC1) are summed. A result of the Monte Carlo simulationis shown in the red.

In SGD, we can reject backgrounds which are unlikely
to be celestial gamma rays by requiring the consistencyof the Compton scattering angles, one derived from theCompton kinematics, and the other derived from the geometrical information. The resulting spectrum in four
operational Compton cameras (three in SGD1 and one
in SGD2) after the background subtraction during the
Crab observation is shown in Figure 2 right. The preliminary result of the Monte Carlo simulation after thesame event reconstruction as the observation data is also
shown in the red spectrum. In the simulation of the Crabnebula observation, we assumed a power-law spectrumwith a photon index of 2.1. The 40–1000 keV ﬂux was
2.0 × 10−8 erg/cm2 /s. During the Crab observation, we
have obtained about 3500 Compton-reconstructed eventsin the range of 50–300 keV, which can also be used forthe polarization measurement, and the careful analysis
for the polarization measurement is in progress.
The Detail of in-ﬂight performance of SGD is summarized in Watanabe et al. (2016), and also the performance of BGO Anti-coincidence detectors are reported
in Ohno et al. (2017) in this proceeding.
Watanabe et al. 2014, NIM-A 765, 192
Watanabe et al. 2016, Proc. SPIE 9905, 13

299

Development and in-orbit performance of all-sky monitoring function of
BGO active shield of the soft gamma-ray detector onboard Hitomi
Masanori Ohno1 , Yasushi Fukazawa1 , Tsunefumi Mizuno1 , Hiromitsu Takahashi1 , Takao Kitaguti1
Yasuyuki Tanaka1 , Jun’ichiro Katsuta1 , Takafumi Kawano1 , Hiroyuki Shirakawa1 , Ikumi Edahiro1 ,Sho Habata1 , Chiho Okada1 , Norie Ohashi1 , Koji Tanaka1 , Tadayuki Takahashi2 , Motohide Kokubun2 ,Shin Watanabe2 , Goro Sato2 , Masayuki Ohta2 , Koichi Hagino2 , Rie Sato2 , Kunishiro Mori2 , Yusuke Uchida2 ,Miho Katsuragawa2 , Ryota Tomaru2 , Hiroki Yoneda2 , Kazuhiro Nakazawa3 , Shunsuke Torii3 , Soki Sakurai3 ,Makoto Sasano3 , Hidetoshi Nishida3 , Shogo Kobayashi3 , Hiroaki Murakami3 , Ko Ono3 , Yuichi Kato3 ,Kastuma Miyake3 , Yoshihiro Furuta3 , Yuki Murota3 , Hiroyasu Tajima4 , Masaomi Kinoshita4 , Kazutaka Yamaoka5,6 ,Katsuhiro Hayashi6, 2, Jun Kataoka7 , Tatsuhiko Saito7 , Masaki Yoshino7 , Takehiro Mimura7 , Yoichi Yatsu8 ,Shin’ya Saito9 , Hideki Uchiyama10 , Takeshi Nakamori11 , Shin’ichiro Takeda12 ,Toshio Nakano13 , Takayuki Yuasa13 , Yuto Ichinohe14 , Daisuke Yonetoku15 , and the Hitomi HXI/SGD team1

Department of Physics, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, Hiroshima 739-8526, Japan2
Department of High Energy Astrophysics, Institute of Space and Astronomical Science (ISAS),
Japan Aerospace Exploration Agency (JAXA), 3-1-1 Yoshinodai, Chuo, Sagamihara 252-5210, Japan3
Department of Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan
4
Solar Terrestrial Environment Laboratory, Nagoya University, Nagoya, Aichi 464-8601, Japan
5
Institute for Space-Earth Environmental Research, Nagoya University,
Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8601, Japan
6
Division of Particle and Astrophysical Science, Graduate School of Science,
Nagoya University, Furo-cho, Chikuka-ku, Nagoya, Aichi 464-8602, Japan
7
Research Institute for Science and Engineering, Waseda University, Shinjuku-ku, Tokyo, 169-8050, Japan
8
Department of Physics, Tokyo Institute for Technology, 2-12-1
Ookayama, Meguro-ku, Tokyo 152-8550, Japan
9
Department of Physics, Rikkyo University, Toshima-ku, Tokyo, 171-8501, Japan
10
Science Education, Faculty of Education, Shizuoka University, Suruga-ku, Shizuoka, 422-8529, Japan
11
Department of Physics, Faculty of Science, Yamagata University,
Kojirakawa, Yamagata 990-8560, Japan
12
Advanced Medical Instrumentation Unit, Okinawa Institute of Science and Technology
Graduate University, Onna, Okinawa 904-0495, Japan
13
Cosmic Radiation Laboratory, Institute of Physical and Chemical Research (RIKEN) Wako, Saitama 351-0198, Japan
14
Department of Physics, Tokyo Metropolitan University, Hachioji, Tokyo 192-0397, Japan
15
Department of Physics, Kanazawa University, Kanazawa, Ishikawa 920-1192, Japan
E-mail(MO): ohno@astro.hiroshima-u.ac.jp
Abstract
BGO active shields of the Soft Gamma-ray Detector (SGD) onboard Hitomi satellite can also act as anall-sky monitor with very large eﬀective area up to soft gamma-ray energy band. We have developed the
hardware and signal processing system of BGO active shield with limited resources. The data processing
system accumulates high- time resolution (16 ms) spectral data (32 channels) once the onboard hardware
trigger is activated for 5.376 sec duration. In addition, the autonomous data accumulation system withmoderate time resolution (4 s) spectral data (128 channels) during about 512 seconds after the trigger
is implemented to cover the long- duration transients. Data accumulation system of BGO active shieldworks well at on-ground calibration tests and in-orbit operation.
Key words: Hitomi: HXI/SGD: Active Shield: Signal Processing

300

Suzaku/WAM
16 ms
64 s
(−8 to 56 s)
50 – 5000 keV
4 ch
400 cm 2

data and folded model

Unit1
Unit2
Unit3

1
0
í1
í2

– Trigger judgement algorithm
Basic idea is following by Suzaku/WAM(Yamaoka et
al. 2017). The trigger is activated when integrated
counts exceed Poisson ﬂuctuation at a certain level. The
trigger time scale can be selected from 1 s or 0.25 s.
The background counts is always accumulating for 8 s
integration time. Energy range for the signal and background can be selected from any ADC channels from 0 to
4095, which is more ﬂexible than that of Suzaku/WAM.
– Quick data transportation
When the trigger is activated, science data is freezed
to the SRAM and the next trigger is disabled until thedata is transferred to DE. To minimize this trigger deadtime, DE is monitoring the trigger ﬂag every 1 s and(dataímodel)/error

2

Energy range
Energy channels
Eﬀective area
(1MeV)

SGD shield
16 ms
5.376 s(−1 to 4.376 s)
150 – 5000 keV
32 ch
800 cm2

0.1

Time resolution
Time coverage

0.01

Table 1. speciﬁcation of all-sky monitoring function of SGD shielddata is transferred after ∼ 10 minutes of the trigger andback to the trigger waiting mode immediately.
– Extra data for longer time coverage
SGD shield also accumulates house keeping (HK) histogram (128 energy channels with 4 sec time resolution)
only during certain time period in the orbit due to thelimitation of the data transfer rate. When the DE ﬁnds
the trigger ﬂag, this HK data allowed to be output for
512 s duration to cover the long time coverage sciencesuch as long GRBs and solar ﬂares.
Until the satellite anomaly on March 26-th, BGO active shield of SGD is operated nominally for 1∼2 weeks.
We had 72 (SGD1) and 44 (SGD2) triggers, but unfortunately, no trigger was generated by real astronomical
event. This is due to non-optimized trigger judgement
condition. However, we conﬁrmed that our developed
data processing system worked as expected for both triggered data and extra longer time duration data. One
astronomical event, GRB 160324A was detected in the
HK data. Figure 1 shows the energy spectrum obtainedby the SGD BGO active shield and the spectrum is well
represented by the single power-law with the photon index of 2.98 (-0.26, +0.29), and the 100−1000 keV ﬂuenceof 1.29 × 10−5 erg cm−2 , which is typical GRB spectral
parameters after the spectral break. Optimized trigger
judgement condition for the trigger energy range can bestudied based on actually obtained in-orbit backgrounddata and we conﬁrmed that if we used optimized trigger condition, GRB 160324A could activate the trigger.
Based on those in-ﬂight performance veriﬁcations, weconclude that all functions of our developed data processing system for all-sky observations worked as expectedin the orbit.

normalized counts sí1 keVí1

All-sky monitering observations in hard X-ray andsoft gamma-ray energy band is very important to understand non-thermal emission mechanism and there aremany all-sky missions such as Swift/BAT, Fermi/GBM,
Suzaku/WAM ..etc. The 6-th Japanese X-ray satellite
Hitomi also has all-sky monitoring function utilizing geometrically large BGO active shield for soft gamma-ray
detector (SGD) (Fukazawa et al. 2014). We have developed dedicated hardwares and data processing system of
BGO active shield of SGD for all-sky observations.
Primary purpose of BGO active shield is to sent theanti-coincident signal to the main detector and thus, theresource for all-sky observation is limited. The science
data of all-sky observations is accumulated in on-board
SRAM only when hardware trigger is activated by transient phenomena and then the stored data is transferredto the digital electronics (DE) and the data recorder.
The SRAM, HLX6228TBR (128 kbytes), similar typeand size of SRAM as Suzaku/WAM is selected to minimize development costs and considering the data sizeof the SRAM, we designed the science data of active
shield of SGD so that we can obtain the spectral datawith high time resolution that can not be obtained by
the Suzaku/WAM. Table1 shows the speciﬁcation of science data of BGO active shield of SGD compared with
Suzaku/WAM. Larger eﬀective area of SGD BGO shieldis advantage for the short transient such as short GRB.
Key development points for SGD BGO active shield aredescribed as following.

200

500
Energy (keV)

1000

Fig. 1. Energy spectra of GRB 160324A obtained by the SGD BGO
active shield. Square, circle and triangle data points show thedata obtained by the diﬀerent detectors and solid lines show the
ﬁtted model.

References
Yamaoka, K. et al., PASJ, submitted
Fukazawa, Y. et al. SPIE, 2014, 9144,2C

301

302

303

Instrumentation

304

305

X-ray Polarimetry Mission PRAXyS
Toru Tamagawa,1
on behalf of the PRAXyS team
1

RIKEN Nishina Center, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
E-mail(TT): tamagawa@riken.jp
Abstract

PRAXyS is a recently proposed NASA’s small explorer mission which aims to open a new window inastrophysics, X-ray polarimetry. In this paper, we are presenting the importance of X-ray polarimetry, the
PRAXyS mission, and performance of ﬂight instrumentation on board PRAXyS.
Key words: X-ray polarimetry — PRAXyS — Time projection chamber

1. Introduction
X-ray photons carry four physical quantities: position(imaging), timing (light curve), energy (spectroscopy),
and polarization. X-ray astrophysics has grown by
mainly using the former three quantities, but not polarization. Although several polarimetry missions were performed, polarization was positively detected from only
one bright and highly polarized source, Crab Nebula,with the OSO-8 satellite (Weisskopf et al. 1976).
Owing to recent technological developments, a highly
sensitive X-ray polarimeter was realized (Costa et al.
2001). Combining with high throughput X-ray mirrors, such X-ray polarimeters allow us to observe faint(∼mCrab) sources with good polarization sensitivity(∼1%) within a realistic observation time (0.1–1 Msec).
The sensitive X-ray polarimetry telescope is a powerful tool to explore the space-time structure and spin of
black holes (Schnittman et al. 2013), vacuum birefringence in the atmosphere of highly magnetized neutronstars (Ghosh et al. 2013), emission mechanism of AGN
jets (Krawczynski et al. 2013) etc.
Polarimeter for Relativistic Astrophysical X-ray
Sources (PRAXyS) led by NASA’s Goddard Space Flight
Center is the highly sensitive X-ray polarimetry satellite, recently proposed as a NASA’s small explorer mission (Jahoda et al. 2016). NASA selected the PRAXyS
project for the Phase A study in July, 2015, together withanother polarimetry mission Imaging X-ray Polarimetry
Explorer (IXPE) led by NASA’s Marshall Space Flight
Center (Weisskopf et al. 2016).
2. PRAXyS mission
Figure 1 shows a schematic view of the PRAXyS satellite. There are two identical X-ray mirrors on top of the
optical bench, and two identical X-ray polarimeters at  
 








  

Fig. 1. A schematic view of the PRAXyS satellite. The satellite
carries two identical X-ray polarimetry telescopes.

those foci. The optical boom will be deployed in space
just after the launch. The satellite will rotate around theoptical axis at the rate of 10 rpm for canceling systematic
uncertainties of the telescopes.
PRAXyS will observe about 30 stellar objects in categories of massive and stellar mass black holes, neutronstars, pulsar wind nebulae, and supernova remnants inthe ﬁrst 9 months after launch with a minimum detectable polarization of down to 1% for a 2 mCrab sourcewith 3.4×106 s observation in the 2–10 keV energy band.
Japanese contribution to the PRAXyS mission was approved as a JAXA’s small project (MoO) in 2016, andNASA’s ﬁnal decision will be made in February 2017 fora launch in 2020.
3. Polarimeter performance
A conceptual design of the X-ray polarimeter on board
PRAXyS is described in ﬁgure 2. An incident X-ray
reacts with a target molecule in the gas cell ﬁlled with
190 Torr dimethyl ether (DME). The diﬀerential cross-306 



 


 


 

Fig. 2. A conceptual drawing of the X-ray polarimeter on board PRAXyS. The time projection chamber technique is employed for imaging
photoelectron tracks. The electron cloud produced along the photoelectron will be ampliﬁed with gas electron multiplier foil (Tamagawa
et al. 2009) and fed into the APV25 digitizing ASIC.

(1 − β cos θ)4

dΩ,

(1)

where φ is the azimuth angle of photoelectron with respect to the X-ray polarization vector, θ is the polar
angle between the incident photon and emitted photoelectron (Heitler 1970). Obviously known from the equation 1, photoelectrons tend to be emitted to the directionof X-ray polarization vector.
To maximise the sensitivity, we have employed a time
projection chamber technique to image the photoelectron tracks (Fig. 2). In this conﬁguration, longer active
volume along the optical axis can be realized, increasing
quantum eﬃciency easily even with a gas detector. We
have eﬃciencies of 0.55 and 0.10 at 3 and 8 keV, respectively, with a 31 cm-long active volume for PRAXyS. The“modulation factor” of the TPC polarimeter is μ=0.52 at
6 keV. Systematic uncertainty is less than 1%. The polarimeter performance derived from NSLS experiments
was summarized in Iwakiri et al. (2016).
The minimum detectable polarization of a polarimeter(99% C.L.) is described as

Pmdp =

4.29√
μ Rs T

1+

Rb
,
Rs

(2)

where Rs and Rb are the source and background countrates, respectively, and T is the observation time. Combining with the mirror eﬀective area, we estimated Pmpd
for PRAXyS as shown in ﬁgure 3. Beside Crab Nebula
we expect positive detections from many astrophysicalsources of several categories at polarization degrees down
to about 1% level.

!

  

  



hν



dσ = r02 α4 Z 5

 mc2  72 4√2 sin2 θ cos2 φsection of photoelectron emission is



!

 

















 









Fig. 3. Minimum detectable polarization Pmdp of PRAXyS for 100 ks
and 1 Ms observations. Theoretically expected polarization degrees from several categories of stellar objects are superposed in
the ﬁgure.

After the conference, the result of NASA’s downselection was announced. PRAXyS was not selected, but
IXPE was selected. A new ﬁeld will be pioneered soon.
References
Black, J. K. et al. 2007, NIMA, 581, 755
Costa, E. et al. 2001, Nature, 411, 662
Krawczynski, H. et al. 2013, arXiv:1303.7158
Ghosh, P. et al. 2013, arXiv:1301.5514
Heitler, W. 1970, The Quantum Theory of Radiation,
Oxford University Press.
Iwakiri, W. B. et al. 2016, NIMA, 838, 89
Jahoda, K. et al. 2016, Proc. SPIE, 9905, 990516
Schnittman, J. et al. 2013, arXiv:1301.1957
Tamagawa, T. et al. 2009, NIMA, 608, 390
Weisskopf, M. et al. 1976, ApJL, 208, 125
Weisskopf, M. et al. 2016, Proc. SPIE, 9905, 990517

307

End-to-End Test of the Mirror and Polarimeter of PRAXyS
Toshio Nakano,1 Toru Tamagawa,1 Wataru Iwakiri,1 Asami Hayato,1 Megu Kubota,1,2 Kazuki Nishida,1,2
Takao Kitaguchi,3 Koyo Toda,3 Noriyuki Tanaka,3 Tsunefumi Mizuno,3 Teruaki Enoto,4 Ikuyuki Mitsuishi,5
Yuzuru Tawara,5 Shota Inoue,6 Kiyoshi Hayashida,6 Keith Jahoda,7 Joe Hill,7 Ross McCurdy,8 and PRAXyS team
1

RIKEN, ishina Center, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
Department of Physics, Tokyo University of Science, 1-3 Kagurazaka, Shinjyuku-ku, Tokyo 162-8601, Japan3
Department of Physical Science, Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima, 739-8526, Japan
4
The Hakubi Center for Advanced Research, Kyoto University, Kyoto 606-8302, Japan
5
Division of Particle and Astrophysical Science, Graduate School of Science,
Nagoya University, Furo-cho, Chikusa-ku, Nagoya, Aichi 464-8602, Japan
6
Department of Earth and Space Science, Graduate School of Science, Osaka University, Toyonaka 560-0043
7
NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
8
Department of Physics and Astronomy, University of Iowa, Iowa City, IA 52242, USA
E-mail(TN): toshio.nakano@riken.jp
2

Abstract
The Polarimeter for Relativistic Astrophysical X-ray Sources (PRAXyS), NASA s small explore mission, will allow us to measure the polarization of the X-ray in the 2–10 keV energy band. We carried out
an end-to-end test of the X-ray optics combining a grazing incidence X-ray mirror and a Time Projection
Chamber Polarimeter for PRAXyS at the NASA Goddard Spacecraft Center 100m X-ray beamline in June
and July 2016. To simulate the spacecraft rotation canceling out instrumental false polarization, the X-ray
tube, mirror and polarimeter in the test were designed to rotate around the beamline axis independently.
The optics was tested with 2.7 keV, 4.5 keV and 6.4 keV X-rays (and bremsstrahlung X-rays) at several
combinations of angles. We checked uniformity of the mirror and performance of the polarimetry.
Key words: X-ray Polarimeter, X-ray mirror, The Polarimeter for Relativistic Astrophysical X-ray
Sources (PRAXyS) mission

1.

Introduction

Although polarimetry in the soft X-ray has been expected to provide a new aspect of astronomy, it is still
untouched for a long time. The PRAXyS mission was
designed to achieve high sensitivity polarimetry in the 210 keV X-ray using two identical telescopes consisting ofa grazing incidence X-ray mirror and a polarimeter (?).
The mirror focuses X-ray onto the polarimeter placedon a focal plane 4.5m away from it. The polarimetercontains Gas Electron Multipliers (GEMs) developed by
RIKEN (?) and 128 read out strips on a 121μm pitch,and is ﬁlled with dimethyl ether at 190 Torr. Using time
projection chamber techniques, the polarimeter images
two-dimentional trajectories of the photoelectrons generated by incident X-ray photons (?). A modulation curveis described as N (φ) = A + B cos2 (φ − φ0 ) , where A
and B are constants, and φ0 is the polarization angleof the source. The modulation amplitude is given as
μ = A/(2A + B).

2. Experimental Conﬁguration
We carried out an end-to-end test of the X-ray optics
consisting of the X-ray mirror and 1/4 engineer model ofthe polarimeter on Summer 2016 at the 100m beamlineof NASA/GSFC. The polarimeter was operated in thesame way as that tested in ?. An electron-impact X-ray
tube was used for the X-ray source, and Rh, Ti and Fe
were used for its metal target. In order to study the instrumental eﬀects for the polarimetry, the X-ray source,
mirror and polarimeter were installed and independently
allowed to rotate around the beam axis.
3. Data analysis and Results
3.1. X-ray Spectra and Modulaction curves
Figure ?? shows spectra of the X-ray from Rh target
taken by the polarimeter. The spectral line at 2.7 keV
is ﬂuorescence line of Rh Lα. The continuum distribution of the bremsstrahlung emission reaches 9 keV whichcorresponds to the voltage of the X-ray tube.
Figure ?? (a) and (b) are modulation curves extracted

308

1

Normalized Counts

0.8

Measured angle (degree) Measured amplitude

Fluorescence line
2.7 keV (Rh L-line)

0.9

Detector angle-45°+45°

0.7
0.6
0.5
0.4

Tube voltage 9kV

0.3
0.2
0.1
0

Bremsstrahlung

1

2

3

4

5
6
Energy (keV)

7

8

9

10

4
3.5
3
2.5
2
1.5
1
0.5
0

Average 2.1

10
5
0
-5
-10
-15
0

20

40
60
80
100
Mirror rotation (degree)

120

140

Fig. 3. The amplitude (Top) and angle (Bottom) of the modulation
curve as functions of the rotational angle of the mirror.

from the ﬂuorescence line and bremsstrahlung continuum, respectively. These curves include the response of
the polarimeter as well as the beam modulation. For
the response correction, we merged X-ray data taken
at the polarimeter angle 45◦ and −45◦ , and obtained
pseudo unpolarized X-ray data. Using this data, wemade response-corrected modulation curves as shown in
Figure ?? (c) and (d). The amplitude of the modulation is small in 2.7 keV as 0.5 ± 0.1%, because ﬂuorescence X-ray photons are unpolarized. On the other
hand, bremsstrahlung photons are generally polarized.
The modulation amplitude of of the bremsstrahlung continuum is 2.0 ± 0.1%. The measured polarization angle
is 1.0 ± 1.4, which consistent with the rotation angle ofthe X-ray tube.

3.3. Source Rotation
To change the polarization angle of the X-ray beam, we
rotated the X-ray tube around the beam axis at -25◦ ,0◦ , and 16◦ , and took the X-ray data with the polarimeter. The polarization angle is expected to be parallelwith the scattering plane of the electrons. Figure ??
plots changes of the measured polarization angles of thebremsstrahlung continuum with the rotational angles of
the X-ray tube. The measured polarization angles correspond to the rotation angles of the X-ray tube.

Moduration curve

1.06(c) Reseponse correctedFluorescence(a) Fluorescence ~2.7 keV

1.04
1.02

1.04

10
0
-10
Slope : 1.10 ± 0.08
Y-intercept: 2.0 ± 1.1

-20

-20

-10

0

10

20

Source rotation (degree)

Red: +45
Blue: -45
Black: Merged

0.96

20

-30
-30

1
0.98

30
Measured angle (degree)

Moduration curve

Fig. 1. Example spectra of the X-rays from the Rh target in the X-ray
tube working at 9kV.

(b) Bremsstrahlung

Fig. 4. Changes of the measured polarization angles with the rotational angles of the X-ray tube.

(d) Reseponse corrected
Bremsstrahlung

1.02
1
0.98
0.96
0.94 <

80 <60 <40 <20 0 20 40
Angle (Degree)60

80<80 <60 <40 <20

0 20 40
Angle (Degree)

60

80

Fig. 2. (Left) Modulation curves made from (a) ﬂuorescence and(b) bremsstrahlung X-rays. Up triangle and down triangle indicate the polarimeter angle 45◦ and −45◦ , respectively. Black
indicates pseudo unpolarized data made by merging the data of
the polarimeter angle -45◦ and -45◦ . (Right) Response-corrected
modulation curves.

3.2. Mirror Rotation
To verify whether that the mirror made any false polarization, we rotated the mirror at 0◦ , 45◦ , 90◦ and 135◦
and ﬁxed the rotational angle of the X-ray tube at 0◦ .
The measured peak angles and amplitudes are constantwithin 1-2 σ errors across the all mirror angles as shownin Figure ??, which means that the systematic errors
introduced by the mirror is small.

4. Summary
We performed an end-to-end test of the mirror and polarimeter for PRAXyS at the 100m X-ray beamline inNASA/GSFC. The measured amplitude was not changed
with the mirror rotation within 1-2 σ errors, whichmeans that the false modulation introduced by the mirror was less than ∼ 1 %. The polarization angles of thebremsstrahlung X-rays were correctly measured acrossthe tested angles of the X-ray tube. Further analysis isongoing for the other data sets.
References
Black, K. et al. 2007, NIMPA, 581, 755
Hill, J. et al. 2016, Proc. of the SPIE, 9905, 99051B
Iwakiri, W. et al. 2016, NINPA, 838, 89
Tamagawa, T. et al. 2009, NIMPA, 608, 390

309

iWF-MAXI : Soft X-ray Transient Monitor on the ISS
N. Kawai,1 Y. Yatsu,1 N. Isobe,1 S. Sugita,1 S. Harita,1 H. Tomida,2 S. Ueno,2 K. Ebisawa,2 T. Dotani,2T. Mihara,3 M. Serino,3 M. Arimoto,4 H. Tsunemi,5 T. Sakamoto,6 T. Kohmura,7 H. Negoro,8
Y. Ueda,9 T. Nakamura,9 T. Tanaka,9 Y. Tsuboi,10 M. Morii,11 M. Yoshida,12 N. Kanda,13 and M. Vagins14
1

Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8550
ISAS/JAXA, 3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa, 252-5210
3
RIKEN, 2-1 Hirosawa, Wako, Saitama 351-0198
4
Waseda University, 3-4-1 Okubo, Shinjuku-ku, Tokyo, 169-8555
5
Osaka University, 1-1 Machikaneyama, Toyonaka, Osaka 560-0043
6
Aoyama Gakuin University, 5-10-1 Fuchinobe, Chuo-ku, Sagamihara, Kanagawa 252-5258
7Tokyo University of Science, 2641 Yamazaki, Noda, Chiba 278-8510
8
Nihon University, 1-8-14 Surugadai, Chiyoda-ku, Tokyo 101-8308
9
Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502
10
Chuo University, Kasuga 1-13-27, Bunkyo-ku, Tokyo 112-8551
11
Institute of Statistical Mathematics, 10-3 Midori-machi, Tachikawa, Tokyo 190-8562
12
Hiroshima University, 1-3-1 Kagamiyama, Higashi-Hiroshima
13
Osaka City University, 3-3-138 Sugimoto, Sumiyoshi-ku, Osaka 558-8585
14
KAVLI IPMU, 5-1-5 Kashiwanoha, Kashiwa, Chiba, 277-8583
E-mail(NK): nkawai@phys.titech.ac.jp
2

Abstract
iWF-MAXI is an X-ray transient monitor mission proposed as an experiment payload on the JEM/EF(Japanese Experiment Module Exposed Facility) of the International Space Station. Its main scientiﬁcgoals are monitor the large part of the sky in X-ray band to ﬁnd and localize astrophysical transient events.
Prompt X-ray emission associated with gravitational wave events is its prime target. Other targets include
gamma-ray bursts, tidal disruption events, stellar ﬂares, nova ignitions, supernova shock breakouts, blackhole/neutron star binaries, and active galactic nuclei (AGN). Its main scientiﬁc instrument is the Soft Xray Large Solid Angle Camera (SLC) with a localization accuracy of 0.1 degrees. Its overview and current
status are presented.
Key words: Instrumentation: all-sky monitor, X-ray transients, gamma-ray bursts, gravitational
waves, X-ray CCD, International Space Station, WF-MAXI

1.

iWF-MAXI and Overview

iWF-MAXI is an X-ray transient monitor mission
proposed as an experiment payload on the JEM/EF(Japanese Experiment Module Exposed Facility) of theInternational Space Station (ISS). It is designed to be theﬁrst X-ray transient monitor in the soft X-ray band (<2keV) with suﬃcient localization capability over most of
the sky with a cadence of 90 minutes. The detector consists of combination of coded mask and CCD in X and Y
coordinates, and is placed on a replaceable small exposed
experiment platform named i-SEEP (Fig.1). iWF-MAXI
has developed with the ﬂight proven technologies fromMAXI (Matsuoka et al. 2009) and Hitomi (Takahashi etal. 2014).
Main scientiﬁc goal of iWF-MAXI is to ﬁnd and local-

ize the X-ray counterparts of gravitational wave events.
The localization enable follow-up observation with X-ray
or optical telescopes, leading to measurements of theirdistance and studies on their environments and progenitors. iWF-MAXI is also expected for detection of othertransients: 1) short soft X-ray transients such as stellar ﬂares, nova ignitions, and supernova shock breakouts, 2) short high-energy transients such as gamma-ray
bursts and tidal disruption events, and 3) Variable X-ray
sources such as black hole binaries, neutron star binaries,
and active galactic nuclei (AGN). Soft X-ray transient
search in the entire sky has huge discovery space for Xray transients.

310cFig. 1. Schematic view of iWF-MAXI on ISS (NASA)

2. Important Role of iWF-MAXI
The ﬁrst gravitational wave event GW150916 was detected by Advanced LIGO (Abbott et al. 2016). The
waveform matched the prediction of merger of a pair ofblack holes. A lot of follow-up observation with electromagnetic telescopes were performed but the FoV ofthem were not enough compared to the localization area
by LIGO because the localization by the gravitational
wave detector alone are so large. No conﬁrmed electromagnetic counterpart detected. In order to associate
the sources with the known astronomical objects, and/or
measure their distances, and identify their origins, detection and localization of electromagnetic counterparts are
sorely needed. Soft X-ray band is one of the promising channel considering the huge energy density at thesource, and yet all-sky monitoring with suﬃcient sensitivity and cadence has never been performed. iWFMAXI will monitor a large area (15% of the entire sky
at any moment) with a high duty cycle (more than 50%of the total real time) to capture rare short X-ray transients. It will localize X-ray transients with positional
accuracy of 0.1 degrees. An alert with the coordinates ofthe source send to the internet for the follow-up observation within several seconds from the in-orbit detection.
With MAXI we have demonstrated the ﬂash alert operation (Negoro et al. 2012).
3. Soft X-ray Large Solid Angle Camera (SLC)
The scientiﬁc instrument of iWF-MAXI is Soft X-ray
Large Solid Angle Camera (SLC). it consists 2 units of
CCD arrays with 1-d coded masks (X and Y directions).
We apply Hamamatsu CCD which is very similar to themodel developed for Hitomi/SXI, and with some changes
in pixel format and packaging (pin grid array instead ofwire bonding). The coded mask for localization is madeof 50 μm thick tungsten that stops X-rays with energies
below 20 keV. The mask element size is 200 μm. The
total opening fraction is about 1/3, determined basedon simulations. The energy range of SLC is 0.7 - 10 keV.
The sensitivity is 100 mCrab in 100 s (5σ) with eﬀective
area of 293 cm2 . Position localization accuracy is 0.1

degrees with FoV of 15% of the entire sky.
We fabricated the bread board model (BBM) of SLC
and performed test operations. Figure 2 shows the BBM
which consists a CCD chamber, a mechanical cooler andelectronics. One CCD was installed inside the chamber.
The set up of the SLC BBM was completed and the basic
operation of the μ-code in the drive electronics has been
conﬁrmed. We plan to evaluate the performance of the
CCDs using this SLC BBM, by the end of this ﬁscal year.
We also performed radiation damage test to the samemodel CCD at HIMAC. After irradiation equivalent to 8
years in orbit, the SLC CCD performance was similar tothat of the spear of the Hitomi CCD (or possibly better),and the charge transfer ineﬃciency was less than 10−4 .
Dark current did not aﬀect the energy resolution anddynamic range of the CCD.

Fig. 2. SLC bread board model

References
Matsuoka, M. et al. 2009, PASJ, 61, 999
Takahashi, T. et al. 2014, Proceedings of the SPIE,
Volume 9144, id. 914425 24 pp. (2014)
Abbott, B.P. et al. (LIGO Scientiﬁc Collaboration andVirgo Collaboration) 2016, Phys. Rev. Lett. 116,061102
Negoro, H. et al. and MAXI Team 2012, Astronomical
Society of the Paciﬁc Conference Series 461, 797

311

A broadband X-ray imaging spectroscopy with high-angular resolution:
the FORCE mission
Koji Mori,1 Takeshi Go Tsuru,2 Kazuhiro Nakazawa,3 Yoshihiro Ueda4 , Takashi Okajima5 ,Hiroshi Murakami6 , Hisamitsu Awaki7 , Hironori Matsumoto8 , Yasushi Fukazawa9 , Hiroshi Tsunemi10 ,and Tadayuki Takahashi11
1

Department of Applied Physics and Electronic Engineering, University of Miyazaki, Miyazaki 889-2192, Japan2
Department of Physics, Kyoto University, Kyoto 606-8502, Japan3
Department of Physics, University of Tokyo, Tokyo 113-0033, Japan
4
Department of Astronomy, Kyoto University, Kyoto 606-8502, Japan
5
NASA/Goddard Space Flight Center, MD 20771, USA
6
Department of Information Science, Faculty of Liberal Arts, Tohoku Gakuin University, Miyagi 981-3193, Japan
7
Department of Physics, Ehime University, Ehime 790-8577, Japan
8
Kobayashi-Maskawa Institute, Nagoya University, Aichi 464-8602, Japan
9
Department of Physical Science, Hiroshima University, Hiroshima 739-8526, Japan10
Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan
11
ISAS/JAXA, Kanagawa 252-5210, Japan
E-mail(KM): mori@astro.miyazaki-u.ac.jp
Abstract
We present our concept of future Japan-lead medium-class mission, FORCE (Focusing On Relativistic
universe and Cosmic Evolution), to be launched in the mid 2020s. FORCE is the direct successor to
the broadband X-ray imaging spectroscopy aspect of Hitomi (ASTRO-H) with signiﬁcantly higher angularresolution. The current design of FORCE deﬁnes energy band pass of 1–80 keV with angular resolution of <15 in half-power diameter, achieving a 10 times higher sensitivity above 10 keV compared to any previous
missions with simultaneous soft X-ray coverage. Our primary scientiﬁc objective is to trace the cosmic
formation history by searching for “missing black holes” in various mass-scales: “buried supermassive
black holes (SMBHs)” residing in the center of galaxies in a cosmological distance, “intermediate-mass
black holes” acting as the possible seeds from which SMBHs grow, and “orphan stellar-mass black holes”
without companion in our Galaxy. In addition to these missing BHs, hunting for the nature of relativistic
particles at various astrophysical shocks is also in our scope, utilizing the broadband X-ray coverage withhigh angular-resolution. FORCE are going to open a new era in these ﬁelds. The satellite is proposed tobe launched with the Epsilon vehicle that is a Japanese current solid-fuel rocket. FORCE carries threeidentical pairs of Super-mirror and wide-band X-ray detector. The focal length is currently planned tobe 10 m. The silicon mirror with multi-layer coating is our primary choice to achieve lightweight, goodangular optics. The detector is a descendant of hard X-ray imager onboard Hitomi (ASTRO-H) replacing
its silicon strip detector with SOI-CMOS silicon pixel detector, allowing an extension of the low energy
threshold down to 1 keV or even less.
Key words: instrumentation: high angular resolution — X-rays: diﬀuse background — accelerationof particles

1.

Scientiﬁc objectives and requirement

Our primary scientiﬁc objective is to trace the cosmic
formation history by searching for “missing black holes”
in various mass-scales. The missing black holes we deﬁne
includes “buried supermassive black holes (SMBHs)”(> 104 M ) residing in the center of galaxies in a cosmo-

logical distance, “intermediate-mass black holes” (102 –104 M ) acting as the possible seeds from which SMBHs
grow, and “orphan stellar-mass black holes” (< 102 M )without companion in our Galaxy. In addition to thesemissing BHs, hunting for the nature of relativistic particles at various astrophysical shocks is also in our scope.

312

In order to achieve these scientiﬁc objectives described
here, a high sensitivity in the hard X-ray band with simultaneous broadband coverage is required. The sensitivity to point sources is generally limited by exposure
time, background, and source confusion, and is getting
better as exposure time increases, but at some point itsaturates due to source confusion limit. Thus, higher
angular resolution is strongly demanded to achieve ahigh sensitivity. Especially, in order to go down to 2–3
×10−15 erg cm−2 s−1 , a requirement from the AGN survey (Mori et al. 2016), the angular resolution of < 15in HPD is necessary.

Angular resolution
Multi-layer Coating
Field of view at 30 keV
Eﬀective Area at 30 keVEnergy range
Energy resolution at 6 keV
Background
Timing resolution
Working temperature

2.

2.1.

Mission design and scientiﬁc instruments

Fig. 1. A schematic view of the FORCE satellite.

FORCE has a weight of about 1 metric ton, and isplanned to be launched with the Epsilon vehicle thatis an ISAS/JAXA solid-fuel rocket. The satellite will be
put into a circular orbit with altitude of 500–600 km andinclination angle of 31 degrees or less, which is similar to
those of previous Japanese X-ray observatories. Figure 1
shows a schematic view of the satellite. FORCE carries
co-aligned, three identical pairs of a supermirror withhigh angular resolution and a focal-plane detector withbroadband response. The supermirror and the detector
are separated by 10 m focal length. Such a long focal
length is necessary to ensure suﬃcient eﬀective area for
hard X-ray focusing and requires an extendable opticalbench (EOB) that can be stowed to ﬁt in the launchfairing and deployed on-orbit. In the current design, the
10 m length is achieved by a combination of 2 m ﬁxed optical bench (FOB) and 8 m EOB, and the mirror moduleis placed at the end of the EOB.
Table 1 summarizes key instrument parameters at thismoment. High angular resolution of < 15 in the broadband of 1–80 keV characterizes this mission. These parameters are deﬁned by the design concepts of the mirror
and detector, which are described in the following sections.

Table 1. Instrument parameters<15 (HPD)Pt/C∼7 ×7 (50% response)
370 cm2
1–80 keV<300 eV (FWHM)
comparable to Hitomi/HXI
several × 10 μs−20±1 ◦ C

X-ray supermirror with high-angular resolution and
Wideband hybrid X-ray detector
The mirror substrates of our X-ray supermirrors are
made based on the single-crystal silicon mirror technology, which has a high potential of making light-weight,
high-angular resolution X-ray optics (Zhang et al. 2016).
In general, higher angular resolution is accompanied withlarger mass and higher production cost. Considering ourlimited resource, we need to ﬁnd out the best compromise among these factors. The silicon mirror can provide
us with a solution for this issue. Although there are still
a number of technical issues to be veriﬁed experimentally toward a ﬂight-ready telescope, the silicon mirror
with multi-layer coating is our primary choice to achieve
light-weight, aﬀordable, high angular optics.
The focal-plane detector of FORCE, wideband hybrid
X-ray imager (WHXI), has the same concept as the hard
X-ray imager (HXI) onboard Hitomi; Si and CdTe hybrid
detector with active shield. Although HXI consists offour layers of double-sided Si strip detectors (DSSD) anda single layer of CdTe double-sided strip (DSD) detector(Nakazawa et al. 2016), WHXI replaces the four DSSD
layers with a single SOI-CMOS pixel detector (SOIPIX)(Tsuru et al. 2014). Low readout noise achievable by
SOIPIX could lower the energy threshold down to 1 keV
and assure a broadband energy response required to the
focal-plane detector of FORCE. CCD is an established Si
detector with low readout noise, but its working temperature, typically < −60 ◦ C, is too low to be placed closely
together with the CdTe DSD, whose working temperature is about −20 ◦ C, in a single camera. SOIPIX has agood time resolution, and what is more, self-trigger function so that anti-coincidence technique can be utilized,
which is necessary to achieve low background.
References
Mori, K. et al. 2016, SPIE, id. 99051O
Zhang, W. et al. 2016, SPIE, id. 99051S
Nakazawa, K. et al. 2016, SPIE, id. 990511
Tsuru, T.G. et al. 2014, SPIE, id. 914412

313

Development of X-ray SOI pixel sensors
for future X-ray wide ﬁeld cameras
Hideki Hayashi,1 T. G. Tsuru,1 T. Tanaka,1 H. Uchida,1 H. Matsumura,1 M. Itou,1 S. Ohmura,1
Y. Arai,2 I. Kurachi,2 K. Mori,3 A. Takeda,3 Y. Nishioka,3 N. Takebayashi,3 S. Yokoyama,3T. Kohmura,4 K. Tamasawa,4 K. Negishi,4 K. Oono,4 K. Yarita,4 S. Nakashima,5
S. Kawahito,6 K. Kagawa,6 K. Yasutomi,6 H. Kamehara,6 S. Shrestha6
1

Department of Physics, Graduate School of Science, Kyoto University,
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan
2
High Energy Accelerator Research Organization (KEK),
1-1 Oho, Tsukuba, Ibaraki 305-0801, Japan3
Department of Applied Physics, Faculty of Engineering, University of Miyazaki,
1-1 Gakuen Kibana-dai Nishi, Miyazaki 889-2192, Japan
4
Department of Physics, Faculty of Science and Technology, Tokyo University of Science,
2641 Yamazaki, Noda, Chiba 278-8510, Japan
5
Institute of Space and Astronautical Science (ISAS)/JAXA,
3-1-1 Yoshinodai, Chuo-ku, Sagamihara, Kanagawa 252-5210, Japan
6
Research Institute of Electronics, Shizuoka University,
3-5-1 Johoku, Naka-ku, Hamamatsu, Shizuoka 432-8011, Japan
E-mail(HH): hayashi@cr.scphys.kyoto-u.ac.jp
Abstract
We have been developing XRPIX for X-ray wide ﬁeld cameras using two-dimensional coded masks.
XRPIX is a monolithic active pixel sensor based on the silicon-on-insulator (SOI) pixel technology. XRPIX
contains a comparator circuit in each pixel. Thus, XRPIX oﬀers a good time resolution better than 10 μs
in addition to the imaging and spectroscopic capabilities comparable to those of X-ray charge-coupleddevices. Recently, we have developed XRPIX5, which has a large imaging area of 21.9 mm × 13.8 mm andthe format of 608 × 384 pixels. We successfully obtained X-ray spectra from almost all regions of XRPIX5
in the frame mode. We found that the readout noise is ∼ 37 e− (rms) and the energy resolution is ∼710 eV (FWHM) at 13.95 keV. Furthermore, we conﬁrmed that the gain of various regions of XRPIX5 is
18.4–18.8 μV/e− .
Key words: monolithic active pixel sensor, SOI pixel technology, X-ray, imaging, spectroscopy

1.

Introduction

X-ray charge-coupled devices (CCDs) are standard imaging spectrometers widely used for focal plane detectors
in modern X-ray astronomy satellites. CCDs have good
performance because of their small pixel sizes (∼ 20 μmsq.) and good energy resolution almost at the Fano
limit (∼ 130 eV in FWHM at 6 keV). However, CCDs
suﬀer from problems such as poor time resolution (afew seconds) and a high non-X-ray background (NXB)
especially above 10 keV due to interactions of cosmic
rays. Therefore, we have been developing XRPIX withhigh-speed readouts and low background. XRPIX is amonolithic active pixel sensor based on the silicon-oninsulator (SOI) pixel technology [1]. Fig. 1 shows the

cross-sectional view of XRPIX. XRPIX contains a comparator circuit in each pixel for hit trigger (timing) andtwo-dimensional hit-pattern (position) outputs in addition to the imaging and spectroscopic capabilities comparable to those of CCDs. The function allows us to readout analog signals only of hit pixels, which is referred toas the event-driven readout mode. Thus, XRPIX oﬀers
a good time resolution better than 10 μs and a highthroughput reaching > 1 kHz. We can also reduce the
NXB by applying the anti-coincidence technique becauseof the good time resolution. In our previous studies,
we successfully demonstrated the X-ray detection by theevent-driven readout [2].

314

BOX (Buried Oxide)(~ 0.2 µm)

#     

"$

241Am
Spectrum !
 

circuity layer(~ 8 µm)

13.95keV

17.75keV






Calibration

Temp: -40 ℃
Vb: 200 V

Pulse Height [ADU]

Si sensor layer(~ 500 µm)

#
Counts/bin



FWHM
~710 eV

20.78keV

Gain: 18.8 µV/e
109Cd







55Fe

26.34keV


















 [ADU]Pulse Height











Energy [keV]

2. Proposal of X-ray wide ﬁeld cameras using XRPIXMAXI has opened up the all-sky monitor and timedomain astronomy in the soft X-ray band below 10 keV.
In order to push out the frontiers, we propose X-ray wide
ﬁeld cameras using XRPIX with two-dimensional codedmasks. The cameras can reduce the NXB and perform
fast timing observation because XRPIX has a much better time resolution than that of CCDs. The angular resolution in principle reaches 2.5 arcmin by adopting thesame pitch size for the coded mask and the distance between of 10 cm XRPIX and the mask. Furthermore, theﬁeld of view reaches 65◦ × 45◦ by adopting the XRPIX
size of 44 mm × 28 mm (This size is about the same sizeof 4 chips of XRPIX5, which is described later.) and thecoded mask size of 88 mm × 56 mm.
3. Evaluation of XRPIX5
Large size chips are essential for the wide ﬁeld cameras.
Recently, we successfully processed XRPIX5, which has
a large imaging area of 21.9 mm × 13.8 mm and theformat of 608 × 384 pixels. The pixel size is 36 μm ×36 μm. XRPIX5 has high resistivity ﬂoating zone wafer(ρ > 2 kΩ cm) with a thickness of 500 μm. XRPIX5 features the largest imaging area in the XRPIX series. In
this chapter, we report the ﬁrst evaluation results of XRPIX5, especially whether the performance changes due
to its large size.
Fig. 2 shows the 241 Am X-ray spectrum and the gain
calibration curve using 241 Am, 109 Cd and 55 Fe X-rays.
We use only single pixel events for the spectrum. We
took the data using the frame mode in which we readout assigned 8 × 8 pixels in each frame. For the present
data, we read out 8 × 8 pixels around [ra, ca] = [304,192] (“ra” means row address, and “ca” means column
address.). We found that the readout noise, the energy
resolution and the gain are ∼ 37 e− (rms), ∼ 710 eV(FWHM) at 13.95 keV and 18.8 μV/e− , respectively.
This performance is nearly the same as XRPIX3b, which
has the best spectral performance in the XRPIX series[3].
We successfully obtained X-ray spectra from the 5 regions shown in Fig. 3. This suggests that we can ob-

Fig. 2. X-ray spectrum of an 241 Am radioisotope and output pulse
height as a function of X-ray energy. The pulse height is shownin analog digital units (ADU). 1 ADU is 244 μV.

[ra, ca]=[456, 288]

[ra, ca]=[456, 96]

XRPIX5Gain : 18.4 µV/e-

Gain : 18.5 µV/enoise : 37 e- (rms)noise : 36 e- (rms)[ra, ca]=[304, 192]
Gain : 18.8 µV/enoise : 37 e- (rms)[ra, ca]=[152, 96]
Gain : 18.7 µV/enoise : 37 e- (rms)raFig. 1. Cross-sectional view of XRPIX.

[ra, ca]=[152, 288]caGain : 18.8 µV/enoise : 38 e- (rms)Fig. 3. Gain of various regions of XRPIX5.

tain X-ray spectra from the entire region of the sensor.
We evaluated diﬀerences in gain of various regions. As
shown in Fig. 3, the gain is found almost uniform (18.4–
18.8 μV/e− ).
4. Summary
We develop XRPIX and propose X-ray wide ﬁeld cameras using XRPIX with two-dimensional coded masks.
In principle, we can realize the small size and fast timing cameras which have angular resolution of 2.5 arcmin
and ﬁeld of view of 65◦ × 45◦ . We successfully processed XRPIX5 with a large imaging area and obtained
X-ray spectra from various parts of XRPIX5. The readout noise is ∼ 37 e− (rms) and the energy resolution is ∼710 eV (FWHM) at 13.95 keV. The gain is found almost
uniform (18.4–18.8 μV/e− ). At present, we conclude
that the performance does not degrade despite the large
chip size. We will further evaluate XRPIX5 including
the veriﬁcation of the event-driven mode.
References[1] T.G.Tsuru et al. 2014, SPIE, 9144, 914412[2] A.Takeda et al. 2013, IEEE, 60, 586[3] A.Takeda et al. 2015, JINST, 10, C06005

315

Micro-satellite for Astrometry and Photometry of Gravitational Wave
Sources
S.Harita,1 T.Ozawa,1 T.Yoshii,1 Y.Yatsu,1 N.Kawai,1 T.Tawara,2 S.Matunaga,2N.Tominaga,3 S.Kulkarni,4 P.Bilgi,4 T.Sakamoto5 and N.Vasquez6
1

Department of Physics, Tokyo Institute of Technology, 2-12-1 Ookayama Meguro Tokyo
Department of Mechano-Aerospace Engineering, Tokyo Institute of Technology, 2-12-1 Ookayama Meguro Tokyo
3
Department of Physics, Faculty of Science and Engineering, Konan University, Okamoto, Kobe
4
Division of Astronomy, Mathematics & Physics, 105-24, California Institute of Technology, Pasadena, CA 91125
5
Department of Physics and Mathematics, Aoyama Gakuin University, 5-10-1 Fuchinobe, Chuo-ku, Sagamihara
6
Department of Phisics, Escuela Politécnica Nacional, Ladron de Guevara E11-253, Quito 170517
E-mail(SH): harita@hp.phys.titech.ac.jp

2

Abstract
We propose a micro-satellite for surveying EM counterparts of gravitational wave sources and the timedomain astronomy in ultraviolet. Several theoretical study predicts NS-NS mergers may produce strong
UV emission prior to the optical and IR emission. If that is the case wide-ﬁeld UV monitor is usefulfor detection and position determination of EM counterparts. For this purpose we designed a φ200mm
reﬂector and a satellite bus for this mission. The expected limiting magnitude is ∼22 magAB in the
200−300 nm band for 300s exposure. Cooperate with the agile satellite bus system, this satellite system
can survey 100 deg2 with a cadence of 2 cycle h−1 . With an assuming of UV emission from free-neutron
decay, more than one NS-NS merger can be detected with the system.
Key words: Ultraviolet, Time-domain astronomy, Gravitational Wave, Micro-satellite

1. Introduction
In 2015, the ﬁrst gravitational wave (GW) event in human history was ﬁnally detected. The astronomers are
now aiming for detecting the electromagnetic counterparts of GW events. The technical barrier to be overcome is the large error circle of the GW telescopes, typically larger than ∼100 deg2 . Therefore the prompt andaccurate position determination is the necessary step towards the multi-messenger astronomy in the GW-era.
2. Target ProﬁleThe ﬁrst GW event seemed to be generated by a binary black-hole that believed not to produce any radiation. Therefore we expect neutron star (NS)-NS mergers
which produce r-nuclei and these radioactive ejecta result in the kilo-novae. The nature of the kilo-novae isstill unclear and numerous theoretical models are proposed. Part of the claimed that free neutrons in theouter edge of ejecta produce ultraviolet photons due tothe beta-decay(Brian D. Metzger et al. 2015).
In this scenario, as the emitting region is exposed theluminosity is should be higher and prior to optical/IR
photons which experienced scatterings and absorption

Table 1. Mission Requirements
ParameterBand.......
Detection limit.......
Time span.......
Survey area.......
Cadence.......
Delay of Alerts.......
Detection rate.......

Requirement200∼300 nm≥22 magAB∼ 1 hour≥ 100 degree2≥ 2 cycle hr−1
≤ 30 min≥ 1 event yr−1

in the optical thick r-nuclei. Fig. 1 shows the simulated
light curves in various colors of a kilo-novae at 200 Mpc
from the earth. The bluer the peak ﬂux is higher, 22
magAB in U-band. The rise time is about 1 hour whichis shorter than that of the typical kilonovae model (noneutron heating model) by factor of 1/10.
If that is the case, ultraviolet which is brighter andfaster than the other colors can be the best way to detectand determine the position of the GW source. Moreoverthe its light curve has crucial information for constraint
the physical state of the condensed matter in neutronstar, i.e. the equation of states, the r-process nucleo
synthesis.

316

Table 2. Speciﬁcations of the Detector System

Parameter
Diameter.......
Focal length.......
Field of view.......
UV Transmittance.......
PSF.......
CCD size.......
pixel number.......
Quantum eﬃciency.......
Dark current.......

Value
200 mm
430 mm17 deg2
70 %15 arcsec (FWHM)
31 mm × 31 mm
2064 × 2064 pix
60 ∼ 80 %@200∼300 nm
8.3 e− s−1 pix−1 @ −30◦ C

Table 3. System Requirements

Fig. 1. Simulated light curves a NS-NS merger from Brian D. Metzger
et al. (2015).

3. Mission Requirements & Design
Based on the target proﬁle shown in the previous section, the mission requirements are summarized in table
1. It is diﬃcult to transfer all the data to the groundinstantly, we must analyze the raw-data and detect transient source on-board. In this case, two images taken at
the diﬀerent time epoch is required to compare the luminosity of the candidates. If we start the observation justafter the GW detection, we must take more than two
images before the maximum. This result in the lower
limit of

Parameter
Required
Attitude stability
15 arcsec / 10s
Slew speed≥1 deg s−1
Visible Area0∼ 40deg
Power supply
96 W
Data rate
30 MB day−1†: angle from the anti-sun direction.

Designed>15arcsec /10s
8 deg s−10∼40 deg†
120 W @ EOL
36 MB/day−1such as the solar array paddles. For this purpose Hibari
has four 400mm x 400mm solar array paddles mounted
on the end of 400mm long arms and these arms are
driven by motors via reducing gears. Remarkably thismethod enables rapid maneuver faster than 10 deg s− 1
whilst maintaining attitude stability higher than 15 arcsec s− 1. In addition, it does not requires additional spaceand mass for the actuator as is for control moment gy2 −1[Num of tiling]×[FoV]×[Cadence] ≥ 200 deg hr .(1) roscopes (CMGs) but only motors and small electrical
power. Therefore this method can be valuable especially
While the payload size is strictly limited by the piggy- for the sophisticated missions on micro-satellites withback regulation of the H-IIA launching vehicle. The critical limitations in power, weight and space.
Thanks to the VSAC which eﬀectively utilizes the limmaximum diameter of the telescope is φ200 mm that
requires the minimum exposure time of 300 s to detect ited resource, the satellite bus is fall within the criterion
22 magAB . Therefore the ﬁeld of view of the telescope of 50×50×50 cm3 . In addition the power consumption
system should be larger than 17 deg2 . To satisfy the re- of the satellite is below 50 W, which is about 50 % of ourquirements and that can be put into the satellite bus, we previous 50 kg-satellite with CMGs. The science mission
also requires real time RF communication for followingemployed the Riccaldi-Honders optical system.
As the ultraviolet imager to be coupled with the tele- up GW events and broadcasting detection alerts to thescope, we employed a back-illuminated CCD developed ground. For this purpose we employ the Iridium netby Caltech for UV survey missions. The CCD has large work. This shrinks the delay of both up/down-links
physical format of 31 × 31 mm2 and high Q.E. ≥ 80%. within 30min on average. With this satellite, more thanThat fulﬁll the our mission requirements. The speciﬁca- one UV counterparts of GW sources and a few tens oftions of the detector systems are summarized in Table core-collapse supernovae and the other unknown phenomena can be detected per year.
2.
4. System Design
The above mission strategy requires both high-speed andquite stable attitude control system. For this mission weare developing a novel attitude control method, namely
Variable Structure Attitude Control; VSAC (Kyosuke
Tawara, and Saburo Matunaga 2015) which utilizes
the anti-torque generated by swinging the appendages,

References
Brian D et al.“Neutron-powered precursors ofkilonovae”,MNRAS, 446, 1115-1120, 2015.
Kyosuke Tawara, and Saburo Matunaga, “On Attitude Control of Microsatellite Using Shape Variable Elements,” The 24th Workshop on JAXA: Astrodynamics
and Flight Mechanics, Sagamihara, July 2015.

317

Development of the Solar Neutron Monitor on the ChubuSat-2 Satellite
Kazutaka Yamaoka,1 Takuya Miyazawa,1 Shinji Oseki,1 Yasunobu Babazaki,1 Yuki Hayashi,1 Daisuke Ishihara, 1
Kazuya Itoh,1 Hidehiro Kaneda,1 Hiroaki Kawahara,1 Shin Kubo,2 Koji Matsushita,1 Kikuko Miyata,1Hosei Nagano,1 Yasutaka Narusawa,1 Masaki Nishino,1 Sosuke Noda,1 Daeil Park,1 Masaaki Sadamoto,2Hiroyasu Tajima,2 Keisuke Tamura,1 Hidetaka Tanaka,1 Dao Ngoc Hanh Tam,1 Toyoki Watabe,1 and Qidong Zhou1
1

2

Nagoya University, Nagoya, Aichi, Japan
Clear Pulse Co., Ltd., Ohta-ku, Tokyo, Japan
E-mail(KY): yamaoka@isee.nagoya-u.ac.jp
Abstract

The solar neutron observation is a key in understanding of the ion accerelation mechanisms at the Sun
surface since neutrons are not aﬀected by the magnetic ﬁeld around the Sun and intersteller mediums.
However, there have only been a few tens of detections so far since its discovery in 1982. ChubuSat
is a series of 50-kg class microsatellites jointly developed by Nagoya university, Daido university, andaerospace industrial companies in the Chubu area of central Japan. ChubuSat-2 was selected as one of thefour piggyback payloads of the X-ray astronomy satellite ASTRO-H in 2014 summer, and was successfullylaunched by the H-IIA launch vehicles from the JAXA Tanegashima Space Center (TNSC) on February
17, 2016. ChubuSat-2 carries two mission instruments, the radiation detector (RD) and the infrared
camera. The main mission of ChubuSat-2 is devoted to monitoring neutrons and gamma-rays which
can be background sources for ASTRO-H celestial observations. The mission also involves solar neutron
observations which were originally proposed by graduate students who join the leadership development
program for space exploration and research, a program for the leading graduate schools (LGS) at Nagoya
University. The RD has a similar detection area and eﬃciency to those of the neuron detector, SEDA-AP
FIBer detector (FIB) on the International Space Station (ISS), but is expected to have a lower backgroundthan that thanks to a much smaller mass of the micro-satellite. In this paper, we will describe developmentof the ChubuSat-2 satellite and the RD, and the current status in orbit.
Key words: solar physics — microsatellite — instrumentation — neutron — gamma-ray

1.

The ChubuSat-2 Satellite

ChubuSat is a series of 50kg-class micro-satellites developed by a collaboration among Nagoya and Daido University, and aerospace industry companies (Monozukuri
Aerospace Support Technology Team: MASTT) aroundthe Chubu region in Japan. This region including
Nagoya city is at the center of Japan, and it is the coreregion of Japanese aerospace industries. The aim of thisproject is to form the basis of industrialization of the
micro-satellite business and revitalize the Chubu region.
Another purpose is to foster global human resources withspecialized knowledge about the aerospace ﬁeld utilizing
development of the ChubuSat satellite.
ChubuSat-2 (Yamaoka et al. 2016) is the second
ChubuSat following ChubuSat-1 launched on November
2014. The primary mission for ChubuSat-2 is to monitor
radiation backgrounds for the ASTRO-H astronomicalsatellite at the same time and orbit. We also added solar neutron observations to the radiation measurements.

Solar neutrons were ﬁrst observed in 1982, but there have
only been about a few tens of successful detections so far.
Neutrons are not aﬀected by the magnetic ﬁeld on the solar surface and in the interstellar medium unlike protons
and electrons, they can be direct probes to clarify the ion
acceleration mechanisms. However, the currently working space detector, SEDA-AP FIB (Muraki et al. 2012)
on the ISS, was aﬀected by the neutron background dueto the ISS itself with a huge mass (∼400 ton). The tiny
ChubuSat-2 is expected to have a much smaller neutron
background than SEDA-AP. Such unique mission was
originally proposed by graduate students who belong to
the Frontier Space Program in the LGSs 1 at Nagoya
University. This educational program aims to developinternational leaders who can spur innovations that will
expand the space utilization using the ChubuSat development.
*1 http://www.frontier.phys.nagoya-u.ac.jp/index.html

318

Fig. 1. A picture of the ChubuSat-2 Satellite.

2. Radiation Detector (RD)The radiation detector (RD) has been newly designed
and developed in collaboration with Clear Pulse Co.,
Ltd. by Nagoya University. The RD can mainly detect both neutrons and gamma-rays among radiations.
It is composed of multi-layered plastic scintillator bars
and a 10×10 GAGG (Ce: Gd3 Al2 Ga3 O12 ) scintillator
array. The plastic scintilators with 10 ×10 layers are utilized for chasing a track of protons recoiled by incident
neutrons. Gamma-rays can be detected in principle ofCompton camera using plastic scintillators as scatterers
and the GAGG array as absorbers. Both scintillators areread out with the novel photo-sensor Multi-Pixel Photon
Counter (MPPC) developed by Hamamatsu Photonics.
The MPPC is very compact with a high quantum efﬁciency and low operation voltage at around +55 V.
Thanks to a use of the MPPCs, we realized very compact detector that can be installed on the micro-satellite.
The size is 15.2 cm × 17.0 cm × 18.5 cm, and the weightis approximately 6.2 kg. The RD is attached to the −Y
panel just behind the solar panel, and always points towards the Sun direction. The speciﬁcation of the RD
is shown in table 1. Figure shows the cosmic-ray muontrack in the RD on the ground. Tracks for simultaneous
events can be identiﬁed in the plastic scintillator bars.

Fig. 3. Muon tracks in the Radiation Detector on the ground. Left:
X-Z plane, Right: Y-Z plane.
Table 1. Speciﬁcations of the radiation detector.

Detector

Geometrical Area
Energy Range
Detection Eﬀ.
WeightSize
Nominal Power
Down-linked Data

100 plastic scintillator bars
+ GAGG (Ce) scintillator
array read out with MPPC
100 cm2
30–120 MeV (Neutrons)
200–1000 keV (Gamma-rays)
1∼2% (Neutrons)
1∼2% (Gamma-rays)
6.2 kg15.2 cm × 17.0 cm × 18.5 cm
12 W∼10 Mbyte per day

3. Development and Current Situation
The RD and the ChubuSat-2 satellite were built in
Nagoya University. Environmental tests such as separation shock, vibration, and thermal vacuum test, havebeen carried out at several facilities. After the ﬁnal qualiﬁcation review by JAXA, it was moved to the launch site
at Tanegashima island.
ChubuSat-2 was successfully launched with ASTROH by the H-IIA rocket #30 on February 17, 2016. The
communication between the satellite and ground station
was established in the amateur radio band soon after thelaunch. However, we had a trouble in the interface between the On-Board Computer (OBC) and Power Control Unit (PCU), so the RD has not been turned on yet.
We will still continue to make an eﬀort for recovery of
the satellite.
References
Yamaoka, K. et al. 2016, Trans. JSASS Aerospace Tech.
Japan, 14, Pf 141
Muraki, Y. et al. 2012, Advances in Astronomy, 2012,
379306

Fig. 2. A picture of Radiation Detector (RD).

319

Development of an in-orbit radiation environment monitor
CUBES onboard a small satellite MIST
Nagomi Uchida,1 Tatsuya Nakaoka,1 Hiromitsu Takahashi,1 Tsunefumi Mizuno,1 Yasushi Fukazawa,1Mark Pearce2,3 and Theodor-Adrian Stana2
1 Department of Physical Science, Hiroshima University, hiroshima 739-8526, Japan

2 KTH Royal Institute of Technology, Department of Physics, 106 91 Stockholm, Sweden.
3 The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova University Centre, 106 91Stockholm, Sweden.

E-mail (NU): uchida@astro.hiroshima-u.ac.jp
ABSTRACT

CUbesat x-ray Background Explorer using Scintillators (CUBES) is jointly being developed by
KTH and Hiroshima University, and the detector is currently planned to be launched in 2018 as one
component of MIST satellite. The aims of CUBES are to study the behavior of three scintillators used
for a future satellite mission (Segmented Polarimeter for High eNergy X-rays; SPHiNX) and theradiation environment in the SPHiNX orbit. For gamma-ray polarization measurements ofGamma-ray Bursts, SPHiNX uses the Compton-based hard X-ray polarimeter which consists ofplastic scintillators as the scatter, and Bi4Ge3O12 (BGO) or Gd3Al2Ga3O12 (GAGG) as the absorber.
Compared to BGO, GAGG is the new material with higher light yield and better energy resolution,
but it has not been used much in space. We evaluated the radiation tolerance of GAGG against 60Co
gamma-rays, and studied the shape dependency of GAGG light yield.
KEY WORDS: Gamma-ray: Gamma-ray Burst — satellite: polarization: scintillator

1.

Gamma-ray Burst and SPHiNX satellite

Gamma-ray Burst (GRB) is the brightest
explosion in the universe. The total emission
energy reaches to 1054 erg and some bursts comefrom the early universe (~12.8 billion light-years).
They are thus important probes to the evolutionof the universe. However, the emission
mechanism and physical environment of GRBs
are not understood well.
Some GRBs have the spectrum consisting of two
components and the Band function cannot
represent the shape [1]. One explanation of suchan additional feature is the thermal emission
from the photosphere. To investigate the origin of
the GRB emission components, gamma-ray
polarization measurement is a powerful tool. If
the distribution of the polarization degree of
GRBs is high (~40%)/low, the emission
mechanism is synchrotron from the jet/thermal
emission from the photosphere close to thecentral fireball.
We are planning to observe gamma-ray

polarization of GRBs, which has been detectedfrom only a few sources, with the Segmented
Polarimeter for High eNergy X-rays (SPHiNX).
SPHiNX satellite aims to be launched in 2021,
and the main detector is the Compton-based hard
X-ray polarimeter (Figure 1.), which consists ofplastic scintillators as the scatter, and Bi4Ge3O12(BGO) or Gd3Al2Ga3O12 (GAGG) as the absorber.
GAGG has high light yield and good energy
resolution, but there is little heritage in space use.
Then, we are investigating whether GAGG issuitable for SPHiNX, since the satellite will beplaced in a low earth orbit and be exposed toprimary cosmic rays, and secondary X-rays,
gamma-rays, and neutrons which cause
activations and generate background signals.
2.

CUBES detector mounted on MIST satellite

To verify the GAGG performance on orbit,
CUbesat x-ray Background Explorer using
Scintillators (CUBES) is jointly being developedby KTH and Hiroshima University. The detector
is planned to launch in ~2018 as one component

320of MIST satellite (Figure 2). MIST is theMIniature STudent satellite developed by KTH
students(https://mistsatellite.wordpress.com).
CUBES will fly the three candidate scintillators
for SPHiNX, which are read out by SiPMs (SiPhotomultipliers), and will study the radiation
environment in the SPHiNX orbit. Below, wereport the test of radiation tolerance of GAGG
against 60Co gamma-rays, and the shapedependency of GAGG light yield.
Plastic scintillator

BGO
or
GAGG

Figure 1. SPHiNX detector

3.

Figure 3. Results of radiation tolerance experiment.

Figure 2. MIST satellite

Radiation tolerance of GAGG

There are already reports about GAGG
activation by protons and radiation tolerance by asmall amount of gamma-rays [2,3]. To evaluate
the radiation tolerance furthermore, we
irradiated 31 TBq 60Co gamma-ray to 15 pieces of
GAGGs twice (~180 krad and ~2.7 Mrad).
Figure 2 shows an example of GAGG energy
spectra of 137Cs before and after irradiation. A
photomultiplier from Hamamatsu (model R7899)
was used for readout. The right sub-figure show acomparison of before, after first, and second
exposure. Additional phosphorescence appears in
the lower energy part, but the energy resolution
is not affected. Right figures show the time
variation after the second irradiation. The
phosphorescence decreases in the order of hours.
The irradiated gamma-rays of several Mradcorrespond to the amount of a few thousand years
in the low earth orbit in space. Therefore, thegamma-ray damage about the phosphorescenceand energy resolution is negligible for GAGGs on
CUBES/SPHiNX.
4.

Shape dependency of GAGG

For the SPHiNX absorber, we are planning to use
flat GAGG scintillators (60 x 27.5 x 5 mm3) tocover the large area. In general, the bigger the
scintillator size, the lower the energy resolution,
when the scintillation light is re-absorbed and/orreflected multiple times at the surfaces. We thus

Figure 4. Results of shape dependency

measured the light yield and energy resolutionsof various dimensions of GAGGs (the surface is 5x 5 mm2, and the length is from 5 to 30 mm). The
S13360-6050CS SiPM was used at thetemperature of 20 degrees.
As shown in Figure 4 (left), the light yield of
137Cs photoabsorption peak decreases as the
GAGG height increases. However, the energy
resolution is almost constant at ~7%. Since the
GRB spectrum has no line features, 7% is goodenough to trace the detector gain with 511 keVannihilation line.
5.

Conclusion

To study the GRB emission mechanism, thegamma-ray polarization measurement is a
powerful tool. We are preparing the detector for
the future SPHiNX satellite mission in 2021. The
CUBES experiment on-board the MIST satellite,
which is currently planned to be launched in2018, is a good opportunity to evaluate thescintillator performances in orbit and measure
the background environment for SPHiNX. In thispaper, we confirmed the high gamma-ray
radiation tolerance and obtained the shapedependency of GAGG scintillator.
References[1]. M. Axelsson et al. 2012, ApJL, 757:L31[2]. M. Sakamoto et al. 2014, Jinst, Vol. 9, P10003[3]. T.Yanagita et al. 2014, Optical Materials, 2016-2019

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323

Some early results from Astrosat
Biswajit Paul1
1

Raman Research InstituteE-mail(BP): bpaul@rri.res.in
Abstract
Astrosat covers a wide energy band in X-rays from 0.2–150.0 keV and it also has two UV telescopes.
We will show some early results from Astrosat on a variety of objects and discuss further science prospects
with Astrosat and a small satellite mission with a Thomson X-ray polarimeter POLIX.

324

The Hard X-ray Modulation Telescope Mission
Shuang-Nan Zhang1
1

Institute of High Energy Physics, China
E-mail(SZ): zhangsn@ihep.ac.cn
Abstract

A dedicated X-ray astronomy satellite, the Hard X-ray Modulation Telescope (HXMT) will be launchedin the spring 2017. HXMT is composed of three sets of collimated instruments in 1–10 keV (about 400
cm2 ), 5–30 keV (about 1000 cm2 ) and 20–250 keV (about 5000 cm2 ) respectively. The anti-coincidence
detector of HXMT can also be used as a sensitive gamma-ray burst detector between 300 keV to 3 MeV.
In this talk, I will describe the instrumentation and scientiﬁc program of HXMT, including the Galactic
plane scanning survey and monitoring as well as pointed observations of X-ray binaries and pulsars.

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331

List of ParticipantsArimoto, Makoto

Waseda University, Japan

Asai, Kazumi

RIKEN, Japan

Axelsson, Magnus

Tokyo Metropolitan University, Japan

Bozzo, Enrico

ISDC, University of Geneva, Switzerland

Corbel, Stephane

University Paris Diderot & CEA Saclay & Observatoire de Paris, France

Debnath, Dipak

Indian Centre for Space Physics, India

Done, Chris

University of Durham, UK

Ebisawa, Ken

ISAS/JAXA, Japan

Enoto, Teruaki

Kyoto University, Japan

Fatima, Siti

Department of Astronomy, Institut Teknologi Bandung, Indonesia

Fujiwara, Taichi

Tokyo Institute of Technology, Japan

Fukazawa, Yasushi

Hiroshima University, Japan

Furusho, Tae

ISAS/JAXA, Japan

Furuya, Kenichiro

University of Miyazaki, Japan

Hagino, Kouichi

ISAS/JAXA, Japan

Hanyu, Chifumi

University of Miyazaki, Japan

Harita, ShoheiTokyo Institute of Technology, Japan

Hayashi, Hideki

Kyoto University, Japan

Hayashida, Kiyoshi

Osaka University, Japan

Hayato, Masami

RIKEN, Japan

Hidaka, Kouki

University of Miyazaki, Japan

Ho, Anna

Calfolnia Institute of Technology, USA

Hori, Takafumi

Kyoto University, Japan

In ’t Zand, Jean

SRON, Netherland

Inoue, Susumu

RIKEN, Japan

Ishisaki, Yoshitaka

Tokyo Metropolitan University, Japan

Islam, Nazma

Raman Research Institute, India

Isobe, Naoki

Tokyo Institute of Technology, Japan

Itoh, Ryosuke

Tokyo Institute of Technology, Japan

Iwakiri, Wataru

RIKEN, Japan

Jaisawal, Gaurava Kumar

Physical Research Laboratory, India

Jana, Arghajit

Indian Centre for Space Physics, India

Jenke, Peter

University of Alabama in Huntsville, USA

Katsuda, Satoru

Chuo University, Japan

Kawai, Nobuyuki

Tokyo Institute of Technology, Japan

Kawakubo, Yuta

Aoyama Gakuin University, Japan

Kawamuro, Taiki

Kyoto University, Japan

Kawase, Tomofumi

Nihon University, Japan

Kennea, Jamie

Pennsylvania State University, USA

Kimura, Mariko

Kyoto University, Japan

Kinugawa, Tomoya

Institute for Cosmic Ray Research, the University of Tokyo, Japan

Kitaoka, Yoshiki

Aoyama Gakuin University, Japan

Koyama, Shu

ISAS/JAXA, Japan

Kubota, Aya

Shibaura Institute of Technology, Japan

332
Kubota, Megu

RIKEN, Japan

Laplace, Eva

University of Tübingen, Germany

Makishima, Kazuo

RIKEN, Japan

Markowitz, Alex

University of California, San Diego, Center for Astrophysics and Space Sciences, USA

Masumitsu, Takahiro

Nihon University, Japan

Matsumoto, Jin

RIKEN, Japan

Matsuoka, Masaru

RIKEN, Japan

Mihara, Tatehiro

RIKEN, Japan

Miyake, Katsuma

The University of Tokyo, Japan

Mizumoto, Misaki

ISAS/JAXA, JapanMori, Koji

University of Miyazaki, Japan

Morii, Mikio

Institute of Statistical Mathematics, Japan

Murata, Naoto

RIKEN, Japan

Nakahira, Satoshi

JAXA/TKSC, Japan

Nakajima, Hiroshi

Osaka University, Japan

Nakajima, Motoki

Nihon University, Japan

Nakamura, Takashi

Kyoto University, Japan

Nakamura, YumikoChuo University, Japan

Nakano, Toshio

RIKEN, Japan

Nakashima, Shinya

ISAS/JAXA, Japan

Nakazawa, Kazuhiro

The University of Tokyo, Japan

Negoro, Hitoshi

Nihon University, Japan

Nishimura, Osamu

National Institute of Technology, Nagano College, Japan

Noda, Hirofumi

Tohoku University, Japan

Oda, Saeko

Kyoto University, Japan

Oda, Sonoe

RIKEN, Japan

Ohashi, Takaya

Tokyo Metropolitan University, Japan

Ohmori, Norisuke

University of Miyazaki, Japan

Ohno, Masanori

Hiroshima University, Japan

Okada, Chiho

Hiroshima University, Japan

Ono, Ko

The University of Tokyo, Japan

Ono, Yuki

Tokyo Institute of Technology, Japan

Paul, Biswajit

Raman Research Institute, India

Pradhan, Pragati

University Of North Bengal & Raman Research Institute, India

Qiao, Erlin

National Astronomical Observatories, CAS, China

Rodes Roca, José Joaquı́n

University of Alicante, Spain

Rodriguez, Jérôme

Astrophysics division, CEA, France

Saisho, Kouta

Tokyo Institute of Technology, Japan

Sakamaki, Ai

Nihon University, Japan

Sakamoto, Takanori

Aoyama Gakuin University, Japan

Sakano, Masa

Wise Babel, UK

Sasaki, Ryo

Chuo University, Japan

Sato, Kosuke

Tokyo University of Science, Japan

Sawada, Makoto

Aoyama Gakuin University, Japan

Sawano, Tatsuya

Kanazawa University, Japan

Seino, Manami

The University of Tokyo, Japan

333
Serino, Motoko

RIKEN, Japan

Shaw, Aarran

University of Alberta, Canada

Shibata, Kazunari

Kwasan and Hida Observatories, Kyoto University, Japan

Shidatsu, Megumi

RIKEN, Japan

Shigeyama, Toshikazu

Research Center for the Early Universe, The University of Tokyo, Japan

Shiki, Kensei

Hiroshima University, Japan

Singer, Leo

NASA Goddard Space Flight Center, USA

Sudo, Masayuki

RIKEN / Shibaura Institute of Technology, Japan

Sugawara, Yasuharu

ISAS/JAXA, Japan

Sugimoto, JuriRIKEN / Rikkyo University, Japan

Sugita, Satoshi

Tokyo Institute of Technology, Japan

Sugizaki, Mutsumi

RIKEN, Japan

Tachibana, Yutaro

Tokyo Institute of Technology, Japan

Takagi, Toshihiro

RIKEN / Nihon University, Japan

Takahashi, Hiromitsu

Hiroshima University, Japan

Takao, Yusuke

RIKEN / Rikkyo University, Japan

Tamagawa, Toru

RIKEN, Japan

Tanaka, Kazuki

Nihon University, Japan

Tanaka, Koji

Hiroshima University, Japan

Tanimoto, Atsushi

Kyoto University, Japan

Terada, Yukikatsu

Saitama University, Japan

Toda, Koyo

Hiroshima University, Japan

Tomida, Hiroshi

ISAS/JAXA, Japan

Torrejón, José MiguelUniversity of Alicante, Spain

Tsuboi, Yohko

Chuo University, Japan

Tsunemi, Hiroshi

Osaka University, Japan

Tsuru, Takeshi

Kyoto University, Japan

Uchida, Nagomi

Hiroshima University, Japan

Uchiyama, Hideki

Shizuoka University, Japan

Ueda, Yoshihiro

Kyoto University, Japan

Urata, Yuji

IANCU, Taiwan

Van der Horst, Alexander

George Washington University, USA

Wang, Sili

Tokyo Institute of Technology, Japan

Yabe, Seiya

Saitama University, Japan

Yabuki, Ken

Chuo University, Japan

Yamada, Yusuke

Aoyama Gakuin University, Japan

Yamamoto, Takayuki

RIKEN, Japan

Yamaoka, Kazutaka

ISEE, Nagoya University, Japan

Yatabe, Fumiaki

RIKEN / Rikkyo University, Japan

Yatsu, Yoichi

Tokyo Institute of Technology, Japan

Yoneyama, Tomokage

Osaka University, Japan

Yoshida, Atsumasa

Aoyama Gakuin University, Japan

Yoshii, Taketoshi

Tokyo Institute of Technology, Japan

Yuan, Weimin

National Astronomical Observatories, CAS, China

Zhang, Shuang-Nan

Institute of High Energy Physics, China

Zhou, Yuanhui

RIKEN, Japan334$XWKRU,QGH[$6$,.D]XPL/$3/$&((YD

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