Prospects, challenges and evolution of AGN modeling in the Astro-H Era
Rikkyo University, Tokyo, October 21-22, 2015

Monte Carlo Approachto Radiation Modeling of AGN
Hirokazu Odaka
ISAS/JAXA

Linking between Theory and Observation
AGN outflow seen from theory and observation
10

PROGA & KALLMAN

Vol. 616

*

5

*

*

‘Model Co

1. Broken PoweContinuumb

2. Galactic Abso
3. Soft X-ray Em−5Residuals (σ)

690

496
REEVES ET AL.
X-ray
spectrum with Suzaku

0

Hydrodynamic simulation

0.5

1

2
Observed Energy (keV)

5

0
−2

4. Iron K Absor

Second line

−6−4

Residuals (σ)

2Second line

6

Proga & Kallman 2004

7

8
9
Rest Energy (keV)

10

11

Figure 2. Upper panel: Suzaku XIS spectra of PDS 456, plotted as a ratio toa simple absorbed power law (Γ = 2.2). Note that XIS 1 is shown in red,
XIS-FI in black. The position of the two soft X-ray emission lines and of the
Fe K-shell absorption are marked with a star. Lower panel: rest-frame ironK-shell absorption in PDS 456 (against the best-fit, broken power-law continuum). Suzaku XIS-FI is shown in black, XMM-Newton EPIC-pn (2001) in
red. The green line represents the expected position of the Fe xxvi 1s–2p line
(6.97 keV), while the blue lines represent the centroids of the detected absorption
lines in XIS.
(A color version of this figure is available in the online journal.)

5. Iron K Emiss

Reeves et al. 2009

No direct link between them!

Fit statisticse

Notes.
a Improvement ib Broken powephoton indices;
normalization icontinuum fluxc Local Galacticd Emission/Abso
in units photons

We can not directly observe the local properties (density, temperature, etc.).
(∆χ = 12 for one parameter) than for the fit with two
To link theory to observational information, we
need radiation modeling.
absorption lines, while the line width is substantially broadened2(σ = 0.40 ± 0.12 keV or a FWHM of ∼ 30000 km s−1 ).
There is also evidence for broadened ionized iron Kα emis-

Structure of Accreting System•

For radiation modeling, we start with a basic picture of an accretingsystem.

•This structure is common both to stellar mass systems and AGN.

Central Engine- accretion flow onto a
black hole or neutron star- emitting strong radiation

Photoionized Plasma- ionized by strong
radiation- reprocessing intolines and continuum

Release of gravitationalenergy- surroundings provide
matter for accretion.
- some part is outflowing.

Relativistic Jet
Energy release via
mechanical power

Cold Matter- reflecting X-rays

Analogy to Stellar Mass Systems•

Stellar mass systems are useful for a benchmark of AGN modeling.

•

X-ray binaries (black holes or neutron stars)•

Accreting white dwarfs are also helpful since they have comparable
disk temperature.
A&A 529, A120 (2011)

epth for very-high energy (>100 GeV) gamma
4; Protheroe & Stanev 1987; Moskalenko et al.
k 2010). Some Wolf-Rayet stars present an exays with luminosities up to ∼1033 erg s−1 , parogen-rich stars (e.g. Skinner et al. 2010). Thisons would be negligible in Cygnus X-3 whereosity is about 1038 erg s−1 , and will not be conthe accreting compact object is a neutron star,n is emitted by the surface of the star. With atemperature T ns ∼ 106 K (see e.g. Yakovlev &he neutron star provides ∼0.25 keV soft X-rays
GeV photons. The luminosity of the neutron star−1
is too low and will be neglected as well in thedominant source of soft X-rays in Cygnus X-3 isaccretion disk formed around the compact obparts of the disk could be very hot with tempera−107 K and emit copious ∼1 keV soft X-rays inprovide in Sect. 2 and in Appendix A a complete
orption of GeV photons in the radiation field ofk.
observations in X-rays revealed that Cygnus X-3
ex spectral and temporal features. Szostek et al.
d the X-ray spectral states into five distinct
he “hard state” to the “soft state”. In the hardz

Cyg X-3

jetobserverGeVγ< 10 keV − GeV (?)
corona1 keVaccretion
disk

20 eV

ψ

γ

r
φs γ dcounterjetx
Wolf−Rayetstar

Cerutti et al. 2011

Fig. 1. Diagram of the microquasar Cygnus X-3 with the diﬀerent
sources of photons considered in this study. The observer sees the disk

X-ray View of an Accreting Object
Cyg X-3 in a wide band (RXTE)L136

Cyg X-3HIGH-RESOLUTION
in high resolution
SPECTROSCOPY (Chandra-HETG)
WITH CHANDRA

Vol. 53

Disk
Comptonization
Non thermal

Emission lines from stellar wind

Fig. 1.—The 1–10 Å spectrum of Cyg X-3 as observed with the HEG (top) and the MEG (bottom) binned in 0.005 Å bins. The positive and negative firorders have been added, and the spectra have been smoothed with a 3 pixel boxcar filter. The labels indicate the positions of various discrete spectral feature
Figure 1. The absorbed spectral shapes of Cyg X-3 from Szostek&“Hea” is the inelegant label for the resonance, intercombination, and forbidden lines in the He-like ions, plotted at the average wavelength for the complex. Hig
Zdziarski (2004). The 42 observations are averaged into five groups,ionization
the features of interest that were not detected have been labeled in parentheses. The horizontal bars indicate the nominal positions of the gaps between t
ACIS chips; the dithering of the spacecraft will broaden the gaps and soften their edges.

Hjalmarsdotter et al. 2009

Paerels et al. 2000

hard state (blue solid line), the intermediate state (cyan long-dashed line),
the very high state (magenta short-dashed line), the soft non-thermal state
for the HEG (MEG) (D. Dewey 2000, private communication).
(green dot–dashed line) and the ultrasoft state (red dotted line).

enlarged scale. The Si xiv and S xvi continua are readily ap
X-ray features emerging from
Comptonization,
disk
black body, disk
The resolution in the Cyg X-3 spectrum can be checked selfparent. The width of these features is a direct measure of th
consistently by analyzing the width of the zero-order image.
electron temperature in the recombining plasma, and a simp
Unfortunately,
thezero-order image
is affectedby pileup.
the ratioof the hard and soft compactnessesℓ /ℓ where ℓ ≡stellar
Lσ/reflection,
photoionized
wind,
etc.
tell
usHowphysicalconditions
eyeball fit to theshapes indicates kT ∼ 50 eV, which is roughlever,enougheventsarriveduring
the
41
ms
CCD
frame
transinagreementwiththe result of model calculations for optical(rm c ), L is the luminosity, σ is the Thomson cross-section, r isfer, forming a streak in the image, that we can construct anthin X-ray–photoionized nebulae (Kallman & McCray 1982of thesizesystem.
Weregionneed
model
to interpretisunbiased
thethe characteristicof the X-ray emitting
and m aone-dimensional
zero-order distribution thisfrom them.spectrum.
A more detailed, fully quantitative analysis of the spectru
e

3

h

s

T

e

T

e

electron mass. The model is weakly dependent on the value ofℓs ,width of this distribution is consistent with the widths ofThenarrow lines in the spectrum of Capella, which indicates thatwill be required to see whether we can also detect the expecte

0

NASA

20

Fe XX

21

12

O VII

22

Ne IX

i

13

6

O VIII RRC

f

r
i

14
15

23
24

7

O VII
RRC

O VIII β

16
17
18
19

26
27
28

Mg XII α

r

8

f

i

9
Ne X βNe IX RRC

Mg XI

10

f

r
O VIII α

Wavelength (Å)

Evans+ (2009)

Figure 2. Co-added Chandra HETG MEG and HEG gratings spectrum of
NGC 1068 from the entire 440-ks data set. Shown are the principal transitions

25

Mg XI βSi KαN VII δ5f

NVII RRCO VII βr

O VII γ

4
Si XIIIO VII δ3
Si XIII 3p–1s

SI Kα

Si XIV αFe XVII

0Counts s−1 Å−1
0.01
0.02
Fe Kα

O VIII γ

Ne X αNe IX βFe XXIINe IX γ2

NVII α11N VII βN VII γ0Chandra←ImageSpectrum→
Counts s−1 Å−1
5×10 −3
0.01

NGC 1068
AGN outflow

Counts s−1 Å−1
2×10 −3
4×10 −3

Evans et al.

X-

4

X-ray View of an Accreting Object

Monte Carlo Approach to Radiative Transfer
Recent improvement in observational instruments allows us to obtainhigh-quality data. Precise comparison between the data andtheoretical models requires careful treatment of X-ray generation in aplasma and radiative transfer.
Monte Carlo is the simplest and the most reliable way to treat
radiative transfer accurately.
The simulation tracks photons by calculating their propagation andinteractions based on Monte Carlo method in a realistic geometry.
to an observer(escaping)emission
the last interaction(E1 , Ω1 , t1 , x1 )cloud

X-ray source
initial condition(E0 , Ω0 , t0 , x0 )Odaka et al. 2011

MONACOWe have developed a new multi-purpose
calculation framework of X-ray radiation
based on Monte Carlo Simulations forobservational study.

MONACO:
Monte Carlo simulation for
Astrophysics and Cosmology
Alphonse Mucha (1897)

The MONACO Framework
MC Simulation
Geometry building (Geant4)
Particle tracking (Geant4)
Physical processes (original)
Output event list

Analysis
Observation (Imaging/spectroscopy)• Geometry• Physical conditions of matter• Initial conditions of photons
for simulation• Initial conditions of photons(Source function)• Observer’s direction• Time of observation

Observed spectra/Images• Distance of the source• Building geometry and tracking particles: Geant4 toolkit library←sophisticated treatment of complicated geometry (e.g. radiation detector simulation)• Physical processes:
original implementation←need to treat plasma, detailed atomic processes of atoms, molecules and ions←need to treat gas motion that makes Doppler effects (shifts and broadening)Physical Processes
Several physical processes important in accretion-powered sources are implemented.

Cold matter

Photoelectric absorption followed by fluorescence andscattering by electrons bound to neutral atoms or molecules
are responsible for generating spectral features of X-ray
reflection- X-ray reflection nebulae (Odaka et al. 2011)- AGN molecular tori (Furui et al. in prep.)

Gas is ionized by strong X-ray radiations in the vicinity ofaccreting objects. Photoionization and photoexcitation ofvarious ion species result in emission lines and absorption lines
Photoionized plasma- Stellar winds in X-ray binaries (Watanabe et al. 2006)- AGN outflows (Hagino et al. 2015, Hagino et al. submitted)- Super soft X-ray sources (bright WDs) (Wada et al. in prep.)

Comptonization

This process plays an important role in cooling a hot accretion
flow through X-ray radiation- Accretion flows onto NS (Odaka et al. 2014)- Accretion flows onto BH (Odaka et al. in prep.)Photoionization
See also Astro-H White Paper “New spectral features” (arXiv:1412.1172v1)•

Stellar winds illuminated by neutron stars or black holes are beststudied X-ray photoionized plasma. e.g. Vela X-1, Cyg X-3•

Watanabe et al. (2006) determined structure of the stellar wind of
Vela X-1 including the wind mass-loss rate based on ChandraHETG specta and detailed Monte Carlo modeling at different
orbital phases.

•

Before MC simulations, we need to calculate ionization structure(ion fractions and electron temperatures), which is determined by- balance between ionization and recombination, and- balance between heating and cooling.

•

We are using XSTAR (Kallman) for this calculation.

in addition
to theradiative(fluorescence).
These two channels of decay areions are not included in the current
version ofcode. decay
For lower
chargein Fig. 3.9.
Since
fluorescence
yields ofi-like ions, we can treat them as neutral
atoms,
whose
implementation
is light elements are very small (e.g.
neutral silicon),the autoionization(Auger eﬀect) can be dominant over the fl
3.3, except for the K-shell photoionization
cross section.
(More accurate
Thus,
these
twosubsection.)
atomic processes
should be properly treated after the photuch low-charged ions are describedin thenext
The K-shell
photoexcitation
L-shellions(1990)(Li-like, Be-like,. . . ). The autoionizatio
n cross sections are given by theandfitting
formulae by ofBand
et al.
ion by the FAC Three-electron
iontransitionin an excitedcomputed
as well the as radiative
rates. state
he electron number Two-electron
and the atomic
number.

Processes in Photoionized Plasma

















Figure 3.9: Schematic diagram of two decay channels of a three-electron ion.



3.5

Resonance-Auger destruction:a photon at a resonance line energy
will be absorbed via resonanceexcitation and autoionization.

Modeling of Thermal and Bulk Comptoni

Comptonization is one of the main processes which generate X-ray radiation
hematic diagram of photon-ion interactions of a two-electron ion.
cretion flows in X-ray binaries or AGNs. Energy sources of Comptonizat

Stellar Wind in Vela X-1
430

Ion fractionofAL.Si XIV
WATANABE ET

WatanabeVol. 651et al. 2006

XSTAR

No. 1, 2006

X-RAY SPECTRAL STUDY OF VELA X-1

433

Fig. 11.—Map of the H-like ion fraction of Si.

for mass-loss rates of 5:0 ; 10!7 , 1:0 ; 10!6 , 1:5 ; 10!6 , and
2:0 ; 10!6 M" yr!1.
4.2. The Ionization Structure
Figure 11 shows maps of the H-like ion fraction of Si fSi H-likeat three different values for the mass-loss rate. At a mass-loss
X
rate of 5:0 ; 10!7 M" yr!1, the concentration of H-like Si ishighest near the midplane of the binary system, as shown in theleft panel of Figure 11. As the mass-loss rate is increased, thedensity everywhere in the wind is also increased. For a fixed
ionization parameter, the distance to the neutron star has to be
reduced, so the region moves closer toward the neutron star, as
shown in the middle and right panels of Figure 11. Multiplying
the maps in Figure 11 by the local density of the stellar winds
gives the maps in Figure 12. These show the absolute abundancesof H-like Si nSi H-like and the region where the X-ray photons of
the emission line are produced.
According to these simulations, a large fraction of emission
lines from highly ionized ions such as H-like Si are producedin the region between the companion star and the neutron star.
This is the region that is obscured by the companion star during

MC simulationMg

eclipse. It turns out that the mass-loss rate is sensitive to the ratio ofline fluxes between phase 0.5 and eclipse because the size of theregion becomes larger when we increase the mass-loss rate. If
a different value for the terminal velocity is assumed, the massloss rate will differ such that the quantity Ṁ# /v1 remains fixed.

Mg XII

Si 4.3.
XIIIEstimate
Si XIVof the Mass-Loss Rate

The simulations are run for the various values of the massloss rate. In order to calculate the observed spectra at different
orbital phases, we designate the direction of the outgoing photon
by the angle ! that the momentum vector makes relative to the lineof centers (see Fig. 13). The phase 0.5 spectrum is constructedfrom the outgoing photons with ! < 10$ , and the eclipse spectrum
is made from photons with 170$ < ! < 180$ . In this calculation,
photons directly coming from the neutron star are neglected since
only the emission lines and scattered components are of interest.
Although many emission lines are resolved in the Chandra
spectra, we use only the Si xiv Ly" line to estimate the mass-loss
rate. This line is one of the brightest lines observed in the data.
Since the energy of this line is sufficiently high, the uncertainties associated with foreground absorption are lower than the low-energy

Fig. 18.—Emission-line spectra obtained with Chandra MEG. The red lines represent the data and the green lines show the Monte Carlo data convolved with MEG
instrument response.

4
3.5
3
2.5
2
1.5
1
0.5
0
0

Vela X-1 (orbital phase = 0.5)
Chandra HETG
Simulation
Fe

Odaka PhD thesis
300

turbulent velocity: 0 km s-1
Be-likeB-likeC-likeN-likeO-like

250

200
150
100
50

2

4

6

8

10

He-like triplet
f+Li-like

4
H-like
3.5
3
2.5
ir
2
L-ions1.5
1
0.5
0
1.7 1.75 1.8 1.85 1.9 1.95 2 2.05 2.1

First attempt to reproduce Si-K complex of
Vela X-1 spectrum obtained by Chandra.
Model parameters are not optimized yet.

0
1.7

1.72

1.74

1.76

1.78 1.8 1.82
Energy [keV]

1.84

1.86

1.88

1.9

turbulent velocity: 100 km s-1
300
Turbulence suppresses resonance-Auger
250
destruction.
200
150
100
50
0
1.7

1.72

1.74

1.76

1.78 1.8 1.82
Energy [keV]

1.84

1.86

1.88

1.9

Monte Carlo simulation is a powerful tool
to simulate eﬀects of turbulence on a line
spectrum.

Powerful winds driven by active galactic nuclei are often thought to provides a very good representation of the data, simultaneously taking
affect the evolution of both supermassive black holes and their host into account both the broad absorption trough at E < 7210 keV in thegalaxies, quenching star formation and explaining the close rela- XIS data and the factor of ,223 excess of flux at E 5 15225 keV observed
tionship between black holes and galaxies1,2. Recent observations of in the PIN data (see Extended Data Fig. 1 and Methods). The column
24
22
large-scale molecular outflows3–8 in ultraluminous infrared galaxies density, ionization and outflow velocity are NH 5 6:4z0:8{1:3 3 10 cm ,21
support this quasar-feedback idea, because they directly trace the gas logj 5 4:11z0:09
cm and vout,X 5 0.255 6 0.011c, respectively{0:04 erg s
9–12
from which stars form. Theoretical models suggest that these out- (see Extended Data Table 1). The wind has a high covering fraction offlows originate as energy-conserving flows driven by fast accretion- CF,X . 0.85, a value that is confirmed if we include the possible associateddisk winds.
• Proposed connections between large-scale molecular emission in the best-fit model (see Supplementary Information). The windoutflows and accretion-disk activity in ultraluminous galaxies were parameters are consistent overall with the mildly relativistic accretionincomplete3–8 because no accretion-disk wind had been detected. Con- disk winds detected in quasars15–19 with a comparable luminosity.
versely, studies of powerful accretion-disk winds have until now focused
Considering the velocity and variability of the wind we estimate a dis13,14
only on X-ray observations of local Seyfert galaxies
and a few tance from the central supermassive black hole in units of Schwarzschild
15–19
higher-redshift quasars. Here we report observations of a power- radii (rS 5 2 GMBH/c2) (where G is the gravitational constant) of r <•ful accretion-disk
wind with a mildly relativistic velocity (a quarter (152900)rS, consistent with accretion-disk scales (see Supplementarythat of light) in the X-ray spectrum of IRAS F1111913257, a nearby Information). Conservatively adopting a radius of 15rS, we estimate a_ out,X < 1.5M[ yr21. Given the very high lumin(redshift 0.189) optically classified type 1 ultraluminous infrared mass outflow rate of M
galaxy hosting a powerful molecular outflow6. The active galactic osity of this active galactic nucleus and the fact that it is accreting at•
nucleus is responsible for about 80 per cent of the emission, with a about its Eddington limit, it is very likely that the wind is radiationquasar-like luminosity6 of 1.5 3 1046 ergs per second. The energetics driven22.
35 almost all parameters constant_Best-fit
spectra of
allP_theobservations
by keepingof these two types of wide-angle outflows is consistent with the energyThe wind’s
momentum
flux
out,X 5 M out,X vout,X < 6 3 10 dyne
9–12
1
z1:7
exceptthe terminal velocity v∞ and outflowing
angle
conserving mechanismthat is the basis of the quasar feedback
corresponds
to 1:3for
P_ rad 5 LAGN/c θmin{0:9 times that of the AGN radiation
If
observer’s
line
of
sight
crosses
multiplevelocity
components
intheoutflow,ain active galactic nuclei that lack powerful radio jets (such jets are (see Supplementary Information). This is consistent with multiple
an alternative way
to drive molecular
outflows).
spectraldrop canbe observed.
20
The mass of the central supermassive black hole in IRAS F111191
0.9
2007
Very broad absorption
feature of an 2011
3257 is estimated to be MBH < 1.6 3 107M[, where M[ is the solar
mass. The resulting Eddington ratio is LAGN/LEdd < 5, where LAGN is
1.4 0.85ultra-fast outflow (Tombesi
et al. 2015)χ2ν=127.9/107
χ2ν=89.7/96
the luminosity of the active galactic nucleus (AGN), indicating that thev~0.1cv~0.1c v~0.2c v~0.3c
Observersource is very likely to be accreting
at
about
its
Eddington
limit.
Inthe
Outflow
0.8
X-ray spectrum it is relatively bright and not strongly affected by neu1.2
tral absorption21. Therefore, it is arguably the best candidate in which
0.75
to study the highly ionized iron K-band absorbers in this class of object.
IRAS F1111913257 was the subject of a long (250-ks) observation
2013a
2013b
2013c
0.7
1.0
from the Suzaku X-ray astronomy satellite obtained by our group in
May 2013.
The spectra of the three separated X-ray imaging spectro0.65
X-raymeter (XIS) detectors onboard Suzaku independently show a prominentχ2ν=103.2/101χ2ν=88.5/89
χ2ν=106.6/95
v~0.1-0.2c
Accretion
diskv~0.1-0.3c
0.8 0.6
absorption feature at the rest-frame energy of ,9 keV (Fig. 1). The
inclusion of a broad Gaussian absorption line in the combined broadband Suzaku XIS and the hard X-ray detector PIN spectrum in the 0.52
0.55
25 keVenergyband provides a very significant improvement of the fit,
Black
hole
0.6
corresponding to a statistical confidence level of 6.5s (see Supplemen0.5
10 14 of PDS 456 by changing only theWe
spectra6 successfully
7 5
8
9 reproduced
10
11 all12 the13tary Information). The absorption line has a rest-frame energy of E 5
Energy(keV)
Rest-frame
energy(keV)terminal velocity and the outflowing angle
z1:00
9:82z0:64
keV,
a
width
of
s
5
1:67
keV
and
an
equivalent
width
E
{0:34
{0:44
2015.08.12
研究会「超巨大ブラックホール降着円盤スペクトルの解釈を巡って」
@ISAS
15
z0:40
1 | Absorptionline
in the Suzaku spectra
spectrum of IRAS
F1111913257.
of {1:31{0:31 keV.
interpreted
as blue-shifted
XXV Hea of
and/or
WeIfapply
ourspectralFemodel
diskFigure
winds
to
the
observedof
1H
between the separated spectra of the Suzaku detectors XIS0 (black),
Fe XXVI Lyman 0707-495,
a resonant absorption
lines, this feature
indicates
an Ratio
andexamine
the
disk
wind
model
explainthis peculiar
structure.
XIS1(red) and can
XIS3 (blue)
and the absorbed
power-law continuum
model inoutflow velocity of about 0.3c (where c is the velocity of light in a vac- the E 5 4212 keV energy range. The ratio for the model including the

AGN Ultra Fast Outflow

Listen to Hagino’s talk!

Some luminous quasars display deep blue-shifted absorption linesof hydrogen/helium-like ions of iron.

Mildly relativistic Disk
(0.3c)wind
windinterpretation
arising from an accretion disk of AGN.

this peculiar
spectralstructurebeluminous
explained
by to theTheCanoutflow
is highly
ionizedby theX-raysfrom
the AGN
Resultsof application
ultra-fast
outflow
central engine. the ultra-fast •outflow?
2×10−3

keV2 (Photons cm−2 s−1 keV−1)

keV2 (Photons cm−2 s−1 keV−1)•

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0

χ

keV2 (Photons cm−2 s−1 keV−1)

2×10−3

−2

•

10−3

2

χ

χ

keV2 (Photons cm−2 s−1 keV−1)

Ratio of spectrum to model

!

10−3

2×10−3

5
Energy (keV)

10

−2

5
Energy (keV)

10Super Soft X-ray Sources (SSS)•

Extremely bright in soft X-rays (below 0.8 keV)•

Thermonuclear burning continuously occurs on a white dwarf
surface, radiating at close to the Eddington luminosity ~ 1038 erg s-1.
-7

-1

•

Very high accretion rate ~ 10 M⦿ yr .

•Possible progenitors of type Ia supernovae.

•

There exists an accretion disk corona (ADC) which displays plentyof emission lines via photoionization.
observerWD

accretion dick corona

accretion disk (optically thick)5. Results (5)τを変えた時のスペクトル変化
Accretion Disk
Corona in(2)SSS
XMM-RGSsimulation

0.1

li
pre

Photons cm−2 s−1 keV−1

normalized counts s−1 keV−1

1

y

ar
n
i
m

χ

0.01
5

0

−5
0.5

0.6
0.7
Energy (keV)

0.8

0.9

10−6 10−5 10−4 10−3 0.01 0.1 1 10 1001000 104 105 106 107

tau=0.2(black),0.1(red),0.05(green),0.025(blue),0.0125(megenta),0.00625(orange)

ADC * SIMPL model (0.4−−0.9 keV)

τ = 0.2

NVII

OVII

OVIII

τ = 0.1 ( 10-2)
τ = 0.05 ( 10-4)
τ = 0.025 ( 10-6)
τ = 0.0125 ( 10-8)
τ = 0.00625 ( 10-10)
0.45

0.5

0.55
Energy (keV)

0.6

0.65

0.7
[10/14]•We attempt to reproduce the high-resolution X-ray spectrum of
CAL 87 (SSS in LMC) in the context of accurate photoionization
modeling.

•

Line red shifts imply that the corona is outflowing.

•

Absence of features by resonance scattering suggests that
optical depth is even thinner than reported by previous studies.
Wada et al. in prep.

Giant Molecular Cloud in the Galactic Center

Giant molecular cloud Sgr B2 hasbeen reflecting a past outburstof SMBH Sgr A*, displaying strong
Fe fluorescence at 6.4 keV.
It shows strong time variability
over a few year scale.
Fig. 1. 6.40 keV maps of the Sgr B2 complex taken with Chandra(upper panel) and Suzaku (lower panel). The coordinates are thegalactic longitude and latitude. Sgr A* is at the right-hand side of thesefigures. The data were smoothed with 1-sigma of 1000 . The Sgr B2
complex is shown by the solid line in the figure (upper panel), while
the Sgr B2 cloud is given by the dashed line (lower panel).

Such objects are good probesof molecular clouds themselves and the black hole activity.

We have to consider multiple scatterings and structure of thecloud. → Monte Carlo simulation

Fi
bla
pa
Ga
co

Physical Processes in Cold Matter
Photoelectric absorption
EPDL97 is used for cross section data.
(Same as G4 Low Energy EM)
Rayleigh
Raman
Compton
Total

Scattering by a bound electron
different from the free electron caseThis process can be divided into threetron binding
makes a great
diﬀerence toprocesschannels
accordingtothe
thefinalparticularlystate offor scattering withall scatteringin otherwords, small recoil energy of the electron. Thescatteredscattering photon energy of 6.4 keV line
theangles,
targetelectron.
bound electrons can be divided into three channels according to the final state oflower
the limit of a(Sunyaev & Churazov 1996)
0.14

tron (Sunyaev & Churazov, 1996). The three channels can be represented0.12
as scattered photon⎧
energy by a free(ground state—Rayleigh scattering)⎨ γ1 + Xi
0.1
electron∗
γ1 + Xi(excited state—Raman scattering)γ + Xi →(2.24)
0.08⎩
γ1 + Xi−1 + e− (free state—Compton scattering)
0.06

e third channel, which is scattering accompanied with ionization, is also called Compscattering. Detailed description of these processes is given in §3.3.2.

3

Photoionized Plasma

0.04
0.02

06.1

6.15

bright astrophysical object that emits UVs or X-rays ionizes surrounding matter via

6.2

6.25
6.3
Energy [keV]

6.35

6.4

6.45

Compton Shoulder by Free Electrons
Free electron temperature

0 eV

2 eV

Scattered photons

5 eV

Scattered photons

Scattered photons

105

105

104

104

104

103

103

103

102

102

102

Counts105

no edges

6.1
6.2 6.3 6.4
6.1 6.2 6.3 6.4
6.1 6.2 6.3 6.4
6.1 6.15 6.2 6.256.3 6.35 6.4 6.45 6.1 6.15 6.2 6.256.3 6.35 6.4 6.45 6.1 6.15 6.2 6.256.3 6.35 6.4 6.45
Energy (keV)

Energy (keV)

Energy (keV)

diﬀuse component, and σabs (E) is the photoelectric absorption cross secti
per hydrogen. The photon index γ of the initial spectrum is 1.8, which isflares of Sgr A* or X-ray emission from Seyfert galaxies.

Simulation Results of the Molecular Cloud•

Giant molecular cloud Sgr B2 has been reflecting
a past outburst of SMBH Sgr A*, displaying Fe
fluorescence line at 6.4 keV.





•


It shows strong time variability over a few year
scale.

Sgr A*



100 pc

•



Such objects are good probes of molecular
clouds themselves and the black hole activity.




Iron 
line (6.4keV) 

Earth

Figure 8.1: Geometricalsetup of the simulation. We assumed three diﬀere
Hard
X-rayskeV)
the SgrB2 cloud,fixing the(20-60
projected
distance


 from Sgr A* to 100 pc.

behind 
+100 pc
center0 pc















129

















2000

2005



before 
-100 pc
2010 2015 2000 2005 2010 2015
Figure 8.4: Time evolution of the morphology of the iron line (left panels) and the hard
Odaka et al. 2011
X-ray (right panels) for Model 2 from the brightest moment at intervals of five years.

Compton Shoulders
Compton shoulder, which is produced by Compton down-scattering ofstrong fluorescence lines (e.g. Fe Kα), has a potential diagnostic power.
hard X-ray

behind (+100pc)reflection from Sgr B2
X-ray(200 ks simulation)centerpc) pc)behind(0(+100-610
beforecenter(-100(0 pc)pc)
1
10-7
before (-100pc)
Comptoncounts s-1 keV-1

photon flux
photon flux

10-6

Fe Kα

-7

10
10-8

-9

10

10-8

10-1

Suzaku CCD
10-2
3

ASTRO-H
SXS
6

-3

10

4
5
Energy [keV]

5

6Energy [keV]
6
Energy [keV]

7

8

9

Comptonshoulder

10

7

7

8

Time Evolution of Compton Shoulder
Response to a short flareradius: R
Time [R/c](0.0-0.1) R/c(0.2-0.3) R/c(0.4-0.5) R/c(0.6-0.7) R/c(0.8-0.9) R/c(1.0-1.2) R/c(1.2-1.4) R/c…

•

hard X-ray

Fe KαComptonshoulderThe shoulder profile can be a probe of the temporal and spatial
properties of the illuminating flare and reflecting cloud.
Odaka, Yoneda et al. in prep.

torus for a small f . On the other hand, the shape of Com
diﬀerent f . We also show the case of the smooth torus fsame. Quantitatively as shown in figure 15, there is only
the Compton shoulder fraction.

AGN Dusty TorusListen to Tanimoto’s talk!
•as a = 0.005,
We next varied
a clumpradius aR
torus
An obscuring dusty molecular torus is believed
to existaroundan
AGN.

total NH at the torus mid-plane and the volume filling fac
respectively.
Then,torus
the numberof clumps
in the torus va• While most X-ray studies have assumed
smoothmodels,infrared
spectra structure.
at higher energy bands are almost id
observations suggest that the torusreprocessed
has a clumpy
As a result, we do not see any dependence on the clum
Comptontorus
shoulder.
• We build a new spectral model to extract
parameters from X-ray

observations.
f= 0.01
f= 0.02

104

105

reprocessed component

f= 0.03104

f= 0.05

f=

f=

f=

103
Photons keV-1

103
102

102

10

10

1

1

10−1

10−1

10−2

f=

f=

f= Smooth

Photons keV-1

Clumpy?
Smooth?

105

10Energy(keV)

2

10

Furui et al. in prep.

10−2
5.8

Fig. 14.— Spectra of the reprocessed component in theComptonization
Optically thickAnalytical/numerical methods are efficient.

Not optically thickComplicated geometry
High energy band

The process is essentially discrete.
→Monte Carlo approach is suitable.

Using Lorentz transformation• determine thenext interactionpoint• see the bulk
motionBulk motion’sframe• select the target
electron• see the thermal
motionpµ
Lorentztransformationp

Target electron’sframe

tau=2.000 kTe=200.000 keV kTbb=0.100 keV• calculate scattering
by a rest electronp µµLorentztransformation

µ
µ
µ
p1
p1
p1
Magnetic field effects can be included.

ν Fν

Lab
frame

Comparison with compPS(Poutanen & Svenson, 1996)

10

−5

10

−6

10

−7

10

−8

10

−9

10

compPSMONACO

− 10

10

−2

10

−1

1

Energy [keV]

10

10

2

10

3

Accretion Flow onto Neutron Stars
Optically thickAnalytical/numerical methods are efficient.

Not optically thickComplicated geometry
High energy band

The process is essentially discrete.
→Monte Carlo approach is suitable.

Accretion column on NS

3

fit NPEX
negative power-law component
positive power-law component

1(Becker & Wolﬀ 2007)
10-1

X-ray

subsonic
flow

photons s-1 cm-2 keV -1

supersonic
flow

X-ray

10-2
10-3
10-4
10-5
10-6

neutron starmagnetic pole

2

3

4

5

6 7 8 9 10Energy [keV]

2030

40 50 60

Odaka et al. 2013, 2014

We found a set of self-consistent solutions that agree with the observationsof accreting NS Vela X-1 and have reasonable model parameters.
column radius = 200 m, kTe = 6 keV, B field=2×1012 G

s.oxfordjournals.org/ at Uchu Kagaku Kenkyujo on January 16, 2014

wards where the hard emission is produced, such lags are naturally
produced (Kotov et al. 2001; Arévalo & Uttley 2006). We note thatthese lags will introduce decoherence, and so suppress variabilityon the time-scale of the lag (Focke, Wai & Swank 2005). However,Cui et al. (2000) show that the time lags in XTE J1550−564 arebelow a few milliseconds at frequencies above 10 Hz so this effectwill be small.
The differences between continuum and fast variability rms spectra are not as great in the harder observations. However, better
signal-to-noise ratio observations of Cyg X-1 reveals similar behaviour in the hard states also (Revnivtsev et al. 1999). Since thetime lags are also present in these states (see, e.g. Miyamoto &Accretion Flow onto Black HolesThe scattering process would be much simpler than the NS case.
anel) and rms No(lower needpanel) spectra
to consider the strong B-field and the bulk motion.
•

A spectrum changes considerably, thes much less alteration.

•

However, the properties of the accretion flow are highly uncertain.

•

We investigate Comptonized radiation of microquasar XTE
J1550–564 on its rise to outburst in 1998.

The source showed strong low frequency QPO, which can be
explained
in thecontext
Axelsson, L.
Hjalmarsdotter
andC. Doneof a precessing inner hot flow.
•

QPOs in power spectral density

Figure 6. Sketch of accretion geometry Axelsson
with different emission
regions
et al. 2013indicated.
results
indicatethat theregion of rapid variability sees very
model (left-hand panels), the PDS of the observation (middle panels) andspectra ofOur
rapidvariability
and QPO(right-handComptonization Model
kTe=14.8 keV, tau=12.8
H

cold diskR

Hot inner flow
We assume cyclo-synchrotron (CS) origin of
the seed photons as well as disk black body.
CS can have a significant role in cooling the
hot inner flow (e.g. Di Matteo et al. 1997).
the flow geometrykTe=14.8 keV, tau=12.8
i=70º±5ºslab geometrykTe=200 keV, tau=0.15
i=70º±8º

l

pre

ry

na
imi

•

Assuming the vertical precession of
the flow (i=70º±5º), we successfully
reproduced the observed level of
QPO. But the first harmonic wassignificantly weaker than observed.

•

This may implies that the flow has
inhomogeneous structure withhigher kTe or nonthermal electrons.
Odaka et al. in prep.

Conclusions
We have developed a general-purpose X-ray
spectrum synthesizer MONACO based onMonte Caro simulation.

Visualization of NS
accretion column

This MC framework will enable us to interprethigh S/N data obtained by modern/future Xray observatories including ASTRO-H.
Collaborators of the MONACO project
Framework: Shin Watanabe, Tadayuki Takahashi
Reflection from molecular clouds: Felix Aharonian, Dmitry Khangulyan
Compton shoulders: Hiroki Yoneda, Andrew Fabian
AGN torus: Shunya Furui, Yasushi Fukazawa, Toshihiro Kawaguchi, Masanori
Ohno, Atsushi Tanimoto, Yoshihiro Ueda
Photoionized stellar winds: Shin Watanabe, Masao Sako
AGN outflows: Kouichi Hagino, Chris Done, Poshak Ghandi
Super soft X-ray sources: Kazuya Wada, Ken Ebisawa, Masahiro Tsujimoto
Accretion onto NS: Dmitry Khangulyan, Yasuyuki Tanaka, Kazuo Makishima
Accretion onto BH: Chris Done
Polarization: Paolo Coppi

Future Plans•

Full simulations of photoionized plasma.
- to calculate ionization structure by using Monte Carlo results.

•

Multi-wavelength, nonthermal processes•

General relativity