ブラックホール降着流の研究~これまでの成果と今後の課題~

大須賀 健 （筑波大）
高橋博之（駒澤大），川島朋尚（東大），野村真理子（呉高専），
北木孝明，恒任優，水本岬希，嶺重慎（京大） ，朝比奈雄太，小川巧未，荻原大樹，井上壮大，内海碧人，高橋幹弥，尾形絵梨花，武者野拓也（筑波大），都丸亮太（Durham Univ.）

降着円盤はなぜ重要？
①宇宙の至る所で自然に形成 ➡ 応用範囲が広い
密度揺らぎ
重力収縮
大部分の質量

球状の天体

大部分の角運動量

周囲の
円盤

円盤銀河

原始惑星系円盤②宇宙の２大エネルギー源のひとつ ➡ 宇宙進化へ影響大恒星

放射

エネルギー源

変換効率

可視光

核反応

0.7%

ブラックホール降着円盤 主にX線 重力エネルギー~10%だからブラックホール天体が明るく光る↗

←E/Mc24

輻射と磁場の重要性
エネルギーの流れ

物質の流れ

輻射エネルギー

輻射力や磁気圧でガスが噴出

磁場/乱流を介したエネルギー変換

重力
エネルギー

ガス

熱エネルギー
降着円盤の断面図

磁場起源の乱流状態＆角運動量輸送輻射と磁場の重要性
1/2 ページ = 左右 80 ミリ × 天地 70 ミリ

磁場の役割①磁気回転不安定 - 角運動量輸送
水平方向への磁力線の引き伸ばし②エネルギー散逸- 円盤加熱

降
着
最 円盤
大
の 理論
再
発 史上
見
！ ,

・真横から見た断面図

回転

角運動量

円盤

A

A

B

ブラック
ホール

B

磁力線

・真上から見た図

角運動量
A
A

B

B

Balbus & Hawley 1991
See also Velikhov 1959

輻射と磁場の重要性
エネルギーの流れ輻射エネルギー

物質の流れ

輻射の役割：
輻射力や磁気圧でガスが噴出①輻射冷却（円盤の温度・厚み）②輻射圧（円盤の厚み）

磁場/乱流を介したエネルギー変換

重力
エネルギー

ガス

熱エネルギー

磁場起源の乱流状態＆角運動量輸送

降着円盤の断面図

輻射と磁場の重要性
1/2 ページ = 左右 80 ミリ × 天地 70 ミリ

輻射加速
磁気加速
1/2 ページ = 左右 80 ミリ × 天地 50 ミリ
輻射力でジェットや円
渦巻き状の磁力線構造が生み
盤風が噴出出す磁気圧でジェット噴出

図 8-6

磁場の力でジェッェットを加速（磁気遠心力もある）

ジェット（ガスの流れ）（電子散乱やライン吸収による輻射力）
光の力でジェットを加速
ジェット

ガス
円盤

光

ブラックホール

円盤

光の力がガスを押す！

磁力線

磁力線のばね効果：
ぐるぐる巻きになるとビヨーンと伸びる

跳ねあげられたガスが飛んでいく

ブラックホール

ガスが円盤表面から飛び出て、ジェットとして飛んで行く

多次元輻射磁気流体力学
荷電粒子は磁力線に引っかかっていることを忘れずに（輻射輸送+磁場+流体+重力）が必須！

磁気流体 vs 輻射輸送
自己重力計算（空間３次元）

輻射の計算(空間３次元＋方向２次元＋振動数)

散乱

光子
吸収

距離だけ
磁気流体計算（空間３次元）
近接作用だけ距離と経路に依存

輻射の計算が最も重い*実際は近似手法を使うが，輻射の計算が難しいことに変わりはない．

人類が観測できるのは，
宇宙から飛来する粒子と重力波を除くと電磁波（光）のみ↓
輻射輸送計算/輻射流体計算は避けて通れない（T_T）*だからこそ理論屋の出番でもある

8

数値計算法については
宇宙物理学の基礎シリーズ第５巻（松本, 須佐, 大須賀 著）
を読んで下さい円盤の多様性Gilfanov et al. 2010

BH天体は様々な輻射スペクトルを示す→ 降着円盤の多様性が起源？

Zdziarski et al. 2010

ADAF/RIAF

標準円盤

スリム円盤

Ichimaru 77; Narayan, Yi 94

Shakura & Sunyaev 73

Abramowicz et al. 88

重力E(少)

重力E(中)

重力E(多)

落下

落下

落下

回転の運動E

回転の運動E

回転の運動E

ガスの熱E

ガスの熱E

ガスの熱E

放射/吸収

放射/吸収

輻射E
L

放射/吸収

輻射E
0.1Ṁ c2L輻射E
0.1Ṁ c2観測者

BH 観測者

L

BH 観測者

BH disk

Mass Accretion Rate

光子捕獲
拡散時間

降着時間

H
r
>
c/τ vr·M = 2πrΣv
τ = σT Σ/mp
2GM
rS = 2c·r
M
H<rS LEdd /c 2 ( r )Mdot > LEdd/c2で発生

0.1Ṁ c2

Three accretion modes
Ohsuga et al. 2009, Ohsuga, Mineshige 2011

RIAF(a) Log(ρ) [g cm-3]

-16

-14

-12

-10

9 -12

Standard

Slim

(b)
) [K]
(a) Log(TLog(ρ)gas[g
cm-3]

-3
(c)
Log(plasma-β)(a)Log(TLog(ρ)) [g(b)[K]cm ]
gas

9.5
-10

10
-8

10.5
-6

-411

-76 -1 -6-58 0 -4(b) Log(Tgas) [K(c) Log(plasma-β)-3
10 1 -2

6 -1

80

8030

80
30

30 80

70

7025

70
25

25 70

60

6020

20
60

20 60

10
40

30

305

530

5 30

20

200

020

0 20

10

10-5

-510

-5 10

0

10

20

30

r/RS

40

0-10

50

4

5

6

7

10
20

15
30

20
40

r/RS
r/R
S
(d) Log(Bϕ2/8π) [erg cm-3]

-100

25
50

5 10

Mass Accretion Rate6(d) Log(Bϕ2/8π) [erg cm-3]
3

5
10

/8π)
6 (e) Log(B78 [erg9cm ] 10
p
2

8

4

-3

5

6

7

8

80

25
70

25

70

20

20

15 30

20 40 25 50

-10 0

2
-3
(f) Log(|B|/B9p) cm10(d)7Log(B]8ϕ /8π)ϕ [erg80

10 0.5 12

114

1.5
16 30

5

10 10

20 15

30 20

r/RSr/RS(f) Log(|Bϕ|/Bp)2(e)
0 Log(B
0.5p /8π) [erg
1

-0.5

6

8

10

12

80
25

70

60

60

10 20

r/R
S
r/R
S
(e) Log(Bp2/8π) [erg cm-3]

-0.5
30 6
80

303
80

9 1

z/RS

4010

08

z/RS

40

Radiation Pressure
15 50
Supported Disk
10 40
& Jet

15
50

z/RS
z/R
S

z/R
z/RS

z/RS

5015
S

Hot rare ed disk
& Magnetized jet

50

Cold thin disk
& Magnetized
wind

7

70
20

60

60

ブラックホール天体の多様性
15

50

40

30
0

30
20

-5

-10

10
セイファート
-10

-10

S

15

20

S

9

25

5

0

50

9.5

-1

10

15

16

17

5

5

10

15

r/RS

20

25

-10
10

0

0

1

S

2

3

4

(h) Log[E0/(e+B2/

0

1

2

5 40

太陽質量の1~10億倍

-10

20

25 80

540

10

15

r/RS
20
30
(i) Log(T
/Trad
)
gas
r/R

(例: M87, SgrA*) 10 50

10
50

-5
20

 

14

10

10

-1

低光度銀河中心核15 60

-5

fi

1

15
60

RIAF

ブラックホール質量

50

13

0

-3

20 70

20

40

0.5

70
20

030

30

SS

0

5

0

50

20

z/RS

z/RS

太陽質量の10倍

r/RS

25

銀河
25
80

15

(低光度状態)

20

-0.5

12

X線連星

10

20

25

0

0

15

r/RS
30
40
(h) Log[E0r/R/(e+B2/8π)]20(g) Log(E ) [erg
cm ]
クェーサー
セイファート30
30Standard disk

40
30

10

10

30

80
70

10

20

8.5

(g) Log(E0) [erg cm-3](高光度状態)
1.8
2
2.2
2.4

1.6

5

10

r/RS
30
40
(g) Log(E) [erg cm-3]
0
r/R

0

50

z/RS

X線連星
r/R

40

10

z/RS

30

z/RS

20

20

20
狭輝線I型
-5

z/RS

10

30
0

-5

超高光度X線源
10

10

50

5

0

20

20

60

40

5

Slim disk (supre-Edd. disk)

30

0

10

z/RS

10

5

50

z/RS

z/R
z/RSSz/RS

z/RS

z/RS

10

40

40

15

50

z/RS

50

光度/ブラックホール質量

15

0 30

-5 20
川島さんの資料を改変

5

10

10

15

20

r/RS
20

30

r/RS

-10 10

25

40

50

0

5

10

10

15

20

r/RS
20

30

r/RS

円盤の共存
三種類の円盤は共存する！？
断面図

降着率

標準円盤

Slim disk
Slim

BH

標準円盤
RIAF
標準円盤

RIAF

詳細はわかっていない

ガス噴出機構の多様性
輻射や磁場がガス噴出を引き起こす．ただし，どの加速が効くか
は円盤のモデルと関連
している

fi

ガス噴出機構の多様性

ブラックホールスピン
BZ効果: 磁場を介してブラック
ホールの回転エネルギーを抽出(c)Kato(Blandford & Znajek 1977)

光度（降着率）magnetic
eld line

ブラックホール質量
ブラックホールスピン

ブラックホール天体を３次元空間で考える時代の到来

ここまでのまとめ• 降着円盤は宇宙の至るところに存在• 理論的には３種のブラックホール降着円盤があり，分岐の原因は降着率• 異なる円盤からは異なるアウトフローが発生するはず• 解析モデルを基本としつつ，シミュ
レーション研究の時代に突入

RIAF（輻射非効率円盤）←ここ*輻射場を解く必要が無いので計算が少し楽
EHTで観測可能な２天体がどちらもRIAFなのは神様からのプレゼントか？

世界初の大局的3DMHD計算- 松元さんおよび松元さんのグループが世界初の３次Magnetic eld

lines
元磁気流体シミュレーションに成功(1999- α粘性の時代から磁気流体計算の時代へ(c)Matsumoto

3DMHDによる円盤構造
円盤の構造はおおまかに解析解と合致

磁気トルクをαパメータに焼き直すと0.01-0.1程度

4fi

Machida et al. 200
see also Hawley & Krolik 2001, Hawley et al. 2001, Kato et al. 2004,

Magnetic eld
lines(c)Matsumoto

局所シミュレーションも活発に行われる．右図は廣瀬らによる輻射
磁気流体計算(c)Hirose

一般相対論的(GR-)MHDへの発展
BZ効果: 磁場を介してブラック
ホールの回転エネルギーを抽出(c)Kato(Blandford & Znajek 1977)

koide et al. 2002,
Sciencemagnetic
eld line

black hole

スピンによるジェット生成を再現した
一般相対論的MHD計算

fifi

局所シミュレーション

BHスピンの効果
Observational Evidence for a Correlation Between Jet Power and Black Hole Spin
低光度円盤＋磁気ジェットのNarayan & McClintock 2012at different radio frequencies ν for four of the five transient BHBs
in our sample. The radio light curves of these four systems were
GR-MHD計算

3

monitored with good time resolution, allowing us to obtain rea-22and open circles correspond respectively to the outbursts studied
by Rodriguez et al. 1995 and Fender et al. 1999). The two lines arefits to the respective data and have a slope of 0.59; writing the spectrum as Sν ∝ ν α , the fit corresponds to α = −0.41. The top rightpanel combines the observations of Hjellming & Rupen (1995) andHannikainen et al. (2000) during an outburst of GRO J1655–40.
The best-fit line corresponds to α = −0.66.2 The lower twopanels show data for XTE J1550–564 (Hannikainen et al. 2009,
α = −0.18) and A0620–00 (Kuulkers et al. 1999). For the lattersource, we do not have enough data points to determine the slope;the line in the plot corresponds to α = −0.4, the average spectral
index of the other three BHBs. In order to enable a fair comparisonof the different objects, we use the fitted lines in the four panels toestimate the peak fluxes (Sν )max at a standard frequency of 5 GHz.
These 5-GHz peak flux values are listed in Table 1.
While each of the above four objects was densely observedin radio during one or more transient outbursts, 4U 1543–47 was
unfortunately not monitored well at radio frequencies during anyof its several outbursts. The only radio data we know of when thesource was bright are those for the 2002 outburst summarized inPark et al. (2004). The strongest radio flux was 0.022 Jy at 1.02675
GHz. Assuming α = −0.4, this gives a flux of 0.0116 Jy at 5 GHz(or only 0.00043 Jy if one corrects for beaming with γjet = 2). We
list this result separately in Table 1 and plot it as a lower limit inFigs 2 and 3 because of the sparse radio coverage. In addition, there
was an anomaly in the 2002 X-ray outburst of this source.
The anomalous behaviour of 4U 1543–47 is apparent by aninspection of figs 4–9 in Remillard & McClintock (2006), whichsummarize in detail the behaviour of six BH transients scrutinized
by RXTE. In panel b of these figures, which displays light curvesof the PCA model flux coded by X-ray state, one sees that only
4U 1543–47 failed to enter the SPL state (green triangles) near the
peak of its outburst, i.e. at the time of the radio coverage reported
by Park et al. Rather, it remained locked in the thermal state (redcrosses) after its rise out of the hard state. This behaviour contrasts
sharply with the behaviour of the other five transients which displayed the strongly-Comptonized SPL state during both the latephase of their rise to maximum and during their early decay phase.
Thus, because of (1) the sparse radio coverage of 4U 1543–47, and(2) the failure of the source to transition out of the jet-quenchedthermal state (Gallo et al. 2003) to the SPL state (which is closely
associated with the launching of ballistic jets), we treat the maximum observed flux of 0.022 Jy as a lower limit. Finally, in sharp
contrast to our finding, we note that figs 5 and 6 in Fender et al.
(2004) indicate a very high jet power for 4U 1543–47. We are unsure how they arrived at their result, but suspect it was based on
infrared data and their equipartition model (see Section 4). If so, an
extension of the present work to infrared data might be worthwhile.
To measure jet power, we scale the 5-GHz peak flux of each
BHB by the square of the distance to the source D. We also divideby the BH mass M since we expect the power to be proportional
to M (this scaling is not important since the range of masses is

BHスピンの効果（BZ効果）で
A. Tchekhovskoy and J. C. McKinney
強力なジェットが噴出する.
A. Tchekhovskoy
and J. C. McKinney
時間変動も激しい.

Jet power

sonably accurate estimates of the peak fluxes. The top left panel
McKinney, Tchekhovskoy,
Blandford (2012)shows data for two separate outbursts of GRS 1915+105 (the solid

BH spin

BHスピンとジェット

Figure 2. Plot of the jet power Pjet as estimated from the maximum radio
flux of ballistic jets (equation 1) vs the measured spin parameter of the BHa∗ for the transient BHBs in our sample. Solid circles correspond to the first
four objects listed in Table 1, which have high quality radio data, and theopen circle corresponds to 4U 1543–47, which has only a lower limit on thejet power. The dashed line corresponds to Pjet ∝ a2∗ , the theoretical scaling
derived by Blandford & Znajek (1977). The data suggest that ballistic jets
derive their power from the spin of the central BH.

パワーが相関するという観測結果！？

only a factor ∼ 2). We thus obtain from the radio observations thefollowing quantity, which we treat as a proxy for the jet power:
Pjet ≡ D2 (νSν )max,5GHz /M.

(1)

It is hard to assess the uncertainty in the estimated values of Pjet .
There is some uncertainty in the values of D and M , but these are
not large. Potentially more serious, the radio flux may not track
jet power accurately. For instance, the properties of the ISM in thevicinity of the BHB may play a role and are likely to vary fromone object to another. Also, the energy released in these roughly
Eddington-limited events will vary (e.g. see GRS 1915+105 in
Fig. 1). Below, we arbitrarily assume that the uncertainty in Pjet
is 0.3 in the log, i.e. a factor of 2 each way.

BHスピンの効果

逆回転

降着率

In the case of GRO J1655–40, the 22-GHz observations did not cover the
ジェットパワー/降着率2peak of the light curve. Hence this point is shown as a lower limit. Similarly,
in A0620–00, the peak was not observed at 0.962 and 1.14 GHz.

3

順回転

JET POWER VS BH SPIN

Fig. 2 shows jet power Pjet plotted against BH spin parameter a∗for the five transient BHBs in our sample. The data are taken from
Table 1. The dashed line has a slope of 2, motivated by the theoretical scaling, Pjet ∝ a2∗ , derived by Blandford & Znajek (1977). The
data points agree remarkably well with this theoretical prediction.
Blandford & Znajek (1977) assumed a slowly spinning BH:a∗ % 1. Tchekhovskoy et al. (2010) obtained a more accurate
theoretical scaling which works up to spins fairly close to unity:
Pjet ∝ Ω2H , where ΩH is the angularfrequency of the BH hori!
zon, ΩH = a∗ (c3 /2GM )/(1 + 1 − a2∗ ). Fig. 3 shows a plot of
Pjet vs ΩH , with the dashed line corresponding to a slope of 2. The
agreement is again very good.

Figure 1. Initial conditions and time-evolution of our fiducial models, A 0.9f (left column) and A0.9f (right column). See Supporting Information for movies.
順回転の場合，降着率は下がるがジェットへのエネルギー変換効[Panels a,f] A vertical slice through the initial conditions. Colour shows fluid frame rest-mass energy density (see colour bar) and thin black lines show levelsof constancy of enclosed
magnetic flux, h i, which represent field lines in the image plane. Magnetic field is axisymmetric, with B' = 0 everywhere and
率が上がる．BHスピンはジェットのパワーと密接に関係がある．
Bz > 0 out to few ⇥ 100rg . Only the magnetic flux enclosed by the thick solid field line eventually falls into the BH. [Panels b,g] Show t- and '-average ofthe magnetically-arrested *ISCOがBHに近いので，効率的にエネルギーが解放できるのかも・・・
state of the simulation. Solid lines show the same contours of h i as in panels (a) and (f). Accretion accumulates so much flux inthe centre that the inner disc is able to push only a fraction of the flux, enclosed by the thick line, into the BH. [Panels c-e,h-j] From top to bottom: rest-mass
Figure1. Initial conditions
time-evolution
our fiducial
A energy
0.9f (leftcolumn)and A0.9f(right BH
column).
See&Supporting
Information
for movies.
energy accretionrate, Ṁc2 ,anddimensionless
BH ofmagnetic
flux,models,
outflow
efficiency,⌘.Tchekhovskoy
Bothand ⌘ saturate
ar t & 6000r
which
BH , andg /c, beyond
McKinney2012
theaccretionflow
ismagneticallyarrested.
Dashedlines
show
time-averages:
aprogradeBH
hasalargerefficiency,
⌘
=
102%,
than
a
retrograde
BH,
⌘
=
34%.
[Panels a,f] A vertical slice through the initial conditions. Colour shows fluid frame rest-mass energy density (see colour bar) and thin black lines show levels

 of constancy of enclosed magnetic flux, h i, which represent field lines in the image plane. Magnetic field is axisymmetric, with B' = 0 everywhere and
Bz > 0 out to few ⇥ 100rg . Only the magnetic flux enclosed by the thick solid field line eventually falls into the BH. [Panels b,g] Show t- and '-average ofvalues,Prm ⇠ 1, thin discs
dothe
notsimulation.
produce centrally-concentrated
content
setup? Tchekhovskoy
al. (2011,
the magnetically-arrested
state of
Solid lines show the same contoursof h ofi asthein initialpanels simulation(a) and (f). Accretion
accumulates soetmuch
flux inmagnetic fields, which suggests that thin discs produce weak or no
TNM11 hereafter) showed that a promising approach is to start withthe centre that the inner disc is able to push only a fraction of the flux, enclosed by the thick line, into the BH. [Panels c-e,h-j] From top to bottom: rest-mass

BHスピンの効果その２
Narayan et al. 2021

逆回転

順回転

最新の研究成果も，順回転の方がジェットへのエネルギー変換効率が高いことを示す．ただし，逆回転であっても無回転BHよりはジェットが強い．← BZ効果
高光度降着円盤の場合は筑波大で計算中（Utsumi et al. in prep.）

BHスピンの効果その３
Fragile et al. 2007

Lense-Thirring precession.
Liska et al. 2019

ジェットは歳差しない！？

観測との協働に向けて
一般相対論的輻射輸送コード

一般相対論的偏光輻射輸送コード

RAIKOUによる計算例

RAIKOU+P（仮）による偏光分布

Tsunetoe et al. 2020

Kawashima et al. 2020

n, A. Sadowski,
˛
R. F. Penna, A. K. Kulkarni

輻射輸送計算によって理論シミュレーションと観測の直接比較→ BHスピンやジェットの加速機構の解明へ*詳細は本日午後の講演でSimulations of Magnetized Advection Dominated Accretion

Figure 1. Poloidal plane of the grid used in the simulations, shown at two zoom levels.

5

未解決問題

6

初期磁場

初期磁場

マルチループ

シングルループ

R. Narayan, A. Sadowski,
˛
R. F. Penna, A. K. Kulkarni

Figure3. Similar
to Fig. 2 but
the ADAF/MAD simulation. The main difference is that here the torus has a single loop of field centered at radius r = 300.
uration of the ADAF/SANE simulation. The top two panels show the mid-plane density and themagnetic
flux threading
the forequatorial
As a result,accretion
causesNotemagnetic
flux of one sign to accumulate around the BH, leading to the MAD state.
adius. Note the extended size of the initial torus, which is required for the extremely long durationof thissimulation.
also thethe magnetic flux, which prevents the accreting gas from reaching the MAD state. The lower two panels show the logarithm of the-magnetic pressure ratio β of the initial torus in the poloidal plane.

対流が顕著

アウトフロー

初期磁場はどこが顕著まで影響があるin the poloidal plane at azimuthal angle φ = 0, and the dashedlines correspond to one density scale height. The main difference
between the two simulations is that the SANE run exhibits moreturbulence compared to the MAD run.
Following Penna et al. (2010), we define the mass accretion
rate Ṁ , the accreted specific energy e, and the accreted specificangular momentum j, at radius r and time t, as follows:! !
Ṁ (r, t) = −ρur dAθφ ,(1)over the course of the simulation. The normalization of the magnetic field is adjusted such that the gas-to-magnetic pressure ratio,β, in the equatorial plane has a minimum value ∼ 100 for each of
the eight loops. Instead of using multiple poloidal loops, anotherway of setting up an ADAF/SANE simulation is to use a toroidal
initial field (e.g., Model A in Igumenshchev et al. 2003 and Model
A0.0BtN10 in McKinney et al. 2012).
The initial magnetic field of the ADAF/MAD run forms a single poloidal loop centered at r = 300 (Fig. 3). The gas accreted by
the BH in this simulation has the same orientation of the poloidal
magnetic field throughout the run, so the accretion flow is maintained in the MAD state. The minimum value of β in the initial
torus is ∼ 50.
The magnetic field construction is described in detail in
Penna et al. (2012).3

θ

のか？

φ

! !
Ė(r, t)1=
Ttr dAθφ ,(2)
Ṁ (r, t)Ṁ (r, t) θ φ
! !
˙ t)
J(r,
1
j(r, t) ==−Tφr dAθφ ,(3)Ṁ (r, t)Ṁ (r, t) θ φ√
where dAθφ = −gdθdφ is an area element in the θ-φ plane,µrrρ is rest mass density, u is the four-velocity, and Tt and Tφ arecomponents of the stress-energy tensor describing the radial flux ofenergy and angular momentum, respectively,

2.3 Preliminary Discussion of the Simulations
The two panels in Fig. 4 show snapshots from the end of the
ADAF/SANE and ADAF/MAD simulation. In each panel, theblack and white streaks and red arrows show velocity streamlines

e(r, t)

=

Ttr

=

Tφr

=

(ρ + Γu + b2 )ur ut − br bt ,(ρ + Γu + b2 )ur uφ − br bφ .

(4)
(5)

Narayan et al. 2012

The quantity u is the internal energy of the gas, Γ is its adiabatic
index which is set to 5/3 in both simulations, and bµ is a four-vector
which describes the fluid frame magnetic field (see Gammie et al.

3

In the notation of Penna et al. (2012), the ADAF/SANE magnetic field
has rstart = 25M , rend = 550M , and λB = 2.5. The ADAF/MAD
magnetic field has rstart = 25M , rend = 810M , and λB = 25.

Figure 4. Left: Snapshot of the ADAF/SANE simulation at t = 200, 000. Black and white streaks as well as red arrows represent flow streamlines. Note theturbulent eddies. The blue dashed lines indicate the density scale height. Right: Snapshot of the ADAF/MAD simulation at t = 100, 000M . There is muchless turbulence.

10
.

M

2003 for details). In equations (1)–(3), the integrals are over the entire sphere (θ = 0 to π, φ = 0 to 2π), and the signs are chosensuch that Ṁ , Ė, J˙ are positive for gas accreting inward. More use-

1

SANE
MAD

未解決問題
電子温度の問題：

フロア値問題：エネルギーを保持しているのは多分陽子だが，光を出すのは主に電子．

ほとんどの降着円盤シミュレーショ
ンでは，回転軸付近で密度フロアがにもかかわらず，MHD計算では電子採用されるため, シミュレーションの温度が不明結果が信用できない．

→２温度計算が必要 Ohmura et al. 2020

→解析モデルを適用

陽子温度

Ogihara et al. 2020

電子温度

*筑波大では小川拓未氏が２温度の輻射MHD計算コードを開発中*MHD近似が成り立たない領域が生じるため, 粒子的な研究も活発に行われている. （冨田さん, 天野さんの講演）

ここまでのまとめ• MHD計算によりRIAFの研究は大幅に発展した.
• BHスピンの効果（特にBZ効果）を解明するた
め, GR-MHD計算が盛んに行われている.

• EHTとGR-MHDの力でM87の解明し, 他のBH天体のメカニズム解明につなげたい.

超臨界降着円盤（Slim円盤）
図 6-2

1/2 ページ = 左右 80 ミリ × 天地 70 ミリ

ブラックホール

電子散乱による輻
射力でジェットや円盤風が発生
スリム円盤

降
着
率

エディントン限界

標準円盤

←ここ

ライアフ

SMBHの急速成長やTDEにも関係

輻射流体シミュレーション
Eddington光度：

4πcGMmp
σT L
GM
=
→
L
=
r2
cmp 4πr 2
σT

光度：~ 3LEdd

天体の最大光度のはずだった・・・
準定常なSuper-Eddington円盤が実現可能であることを世界で初めて実証(Ohsuga et al. 2005)*ただし，磁場を解かずにα粘性を採用

降着率：~ 150LEdd/c2

吸い込まれる光度：~ 10LEdd

SUPER-EDDINGTONはなぜ可能？
Mass density
Radiation

Accretio
Flow

Jets

Strong radiation pressure
supports the thick disk andgenerates the jets, ~0.3-0.7c.

Photons mainly escape throughthe region around the rotation
axis, so that the radiationpressure cannot prevent theaccreting motion.

輻射流体から輻射磁気流体へ
ブラックホールに吸い込まれるガス（円盤の断面図）

ブラックホールn

高橋博之氏（駒澤大学）提供

高橋, 大須賀, 他, 2016

ガス円盤

輻射加速ジェットの発生を確認

ブラックホール

see also Sadowski et al. 2014, Jiang et al. 2014

輻射抵抗

Takeuchi, Ohsuga, Mineshige. 2010
SS433のジェット（光速の約30%）と無矛盾
GRS1915+105のジェット（光速の99%超）は説明できない.
ジェットの細さ（数°以下）も説明できない.

X-RAY SPECTRA
tonpho

w/o Compton

thermal & bulk Compton

thermal Compton

3-

08
07
06

4-

5-

(a) Log(ρ) [g cm-3]

6-7-(b) Log(Tgas) [K]
7

8

9

11

-1
80

70 07

70

70

60 06

60

50 05

50

z/RS

40 04

50

40

40
30
20

01

10 01

10

10

0

0

0

0

04

03

02

01

10

20

30

40

50

10

r/RS

20

30

40

61

(d) Log(Bϕ /8π) [erg cm ]

05

02

8

10

12

14

(e) Log(Bp /8π) [erg cm ]

20

8

10

12

14

16

-0.5

30

40

50

1

1.5

r/RS

Kawashima et al.

(f) Log(|Bϕ|/Bp)0

0.5

Kawashima et al. 2012

When
the mass accretion rate80 is higher, our calculations80in
80 08
NGC1313 X-2
this paper show that SEDs become softer, although our previous
7007
paper70 (Kawashima
et al. 2009)70expected that SEDs become
harder because of Compton upscattering by thermal electrons
60
60
06
1
in the60jet.
Kawashima et al. (2009)
evaluated the photon indicesby calculating
theComptony-parametersinthe
regionτ!
eff 501.
50
50 05
The effective optical depth is calculated by integrating the40
40
40 04 absorption and the scattering
free–freecoefficients from theouter boundary at r = 500rs to the inner boundary at r = 3rs
0.1
03along30 therays with polar angle30θ = constant, for various30θ .
(Here,
we note that the computational
domain for radiation
20
20
20 02
hydrodynamic simulation in Kawashima et al. 2009 was smaller
IC342 X-1
1001 in this work.) Since 10than 10that
Kawashima et al. (2009) did
not take into account the downscattering by cool outflows,
0
0
0 0
10
20 when30
40 accretion
50
10
20
30
40
50
05
04
03
02 they
01 overestimated10
20
30the 40spectral
50
hardening
the
r/RS
r/RS
R/r
r/RS
1
S
rate
is
high.
When
the
mass
accretionrateishigher,agreater]3-mc gre[ )0E(goL )g(
Log(E0) [erg cm-3](h) Log[E0/(e+B2/8π)](i) Log(Tgas/Trad)
number of(g)escaping
photons tends to enter
into the cool outflow
71
61
51
41
31
21
12
13
14
15
16
17
0
1
2
3
-1
0
1
2
3
4
5
surrounding
the
jet
and
they
are
subsequently
downscattered,
80 08
80
80
because the funnel becomes narrower and jet becomes thicker
0.1
70 07
70
for Thomson
scatterings. This 70is the reason why the SEDs
become softer when the mass accretion rate is higher, in
60 06
60
this work.
Nevertheless, the conclusionin Kawashima et 60al.
(2009)
the spectral state changes
from the slim disk one
5
50 0that
50
50
NGC5204 X-1
to the Comptonizing outflows is valid in highly supercritical
vector.
40 04 flows because the SED40of the photons escaping from40
accretion
photosphere, at which the optical thickness measured from the upper boundary is unity. The arrow in panel (g) shows the radiative flux
suchsupercriticalflows
is whose
determinedby exceed
the Comptonand black arrows inpanelhighly(a) indicate
the velocity
vectors
magnitude
the escape velocity. The dashed line in panel (b) is the
3 the gas temperature to the30radiation temperature. All values
30 0of30 are time-averaged over
1 t = 6−7 s. The whiteenergies, (i) and theratio
downscatteringsin the outflow.

X-ray spectra, calculated by the
post-processing radiation transfer,nicely t the observations (ULXs).

Mdot=103LEdd/cviewing angle=10°

hard X-ray is observed sincehigh-energy photons escapethrough the funnel/0E[goL )h(
0
08
07
06
05

S

R/z

R/z

S

hard X-ray is
reduced byCompton down4.3. Comparison with Observationsof ULXs
scattering

Fig. 5.— Two-dimensional distribution of the various quantities forthe gas temperature, (c) the plasma-β, (d) the magnetic energies via t
component, (f) the magnetic pitch, (g) the radiation energy, (h) the renergies, (i) and the ratio of the gas temperature to the radiation temp
and black arrows in panel (a) indicate the velocity vectors whose magnit
photosphere, at which the optical thickness measured from the upper b
vector.
03

z/RS

0

R/r

z/RS

01

04

6

16

R/z

03

66

S

R/z

S

04

8

-3

z/RS

06

01

z/RS

07

21

z/RS

08

41

10

X-ray spectra

2

ϕ

6

0

50

r/RS

The Astrophysical2Journal, 752:18(12pp), 2012 June10-32
] mc gre[ )π8/ B(goL )d(
3-

1

60

20

05

0

Kawashima et al. 2012

30

R/r

01

(c) Log(plasma-β)

10

20 02

8/ B(goL )e(

1

6

80

02

2
p

02

-2

03

S

8

-3

30 03

R/r
01

-4

R/z

01

-5

S

R/z

S

02

-6

80 08

05
04

-7

z/RS

2-

6

z/RS

7

OHSUGA & MINESHIGE

High-energy X-ray
photons are producedby the thermal & bulk
Comptonization.

]3-mc g[ )ρ(goL )a(T(goL )b(
8

8

8

z/RS

M & AGUSHO

swallowed by BH

200LEdd/c
10°200LEdd/c
50°

component, (f) the magnetic pitch, (g) the radiation energy, (h) the ratio of the radiation energy to the sum of the gas and magnetic
the gas temperature, (c)of field, (e) the same but of the poloidal
02 plasma-β, (d) the magnetic
20 the20 energies via the toroidal component
20
Fig. 5.— Two-dimensional distribution of the various quantities for Model A: (a) the density overlaid with the velocity vectors, (b)

02
01

10 01

10

10

S
r/Robservatories
r/RS
Recent 10X-ray(Suzaku,
XMM-Newton,
and 10 20 r/RS30 400.1 50
20
30
40
50
10
20
30
40
50
0
0
0
Chandra)provideus
withprecisespectra
of
ULXs
below
0
0 0
0
0
02
01
05
04
03
02
01
10
20
30
40
50
10
20
30
40
50
10
20
30
40
50
R/r
R/r
10 keV.
In
FigurewecomparetheSEDs
obtained
in
this
work
r/R7,
r/RS
r/RS
S
S
10
10
10
and
the
XMM-Newton
pn
data
of
three
ULXs:
NGC
1313
X-2,
rof seititnauq suoirav Fig.
eht fo5.—
noitTwo-dimensional
ubirtsid lanoisnem
id-owT —.5of .thegiF various quantities for Model A: (a) the density overlaid with the velocity vectors, (b)distribution
0.3 but of the poloidal
1
10
X-1,
are
providedby of field, (e) the same20
t aiv seigrene citengthe
am gaseht temperature,)d( , β-amIC
salp342e20
htthe)c(plasma-β,,erandutareNGCpm
et thes5204
ag magnetic
ehtX-1,20which(c)(d)
energies
viakindly
the toroidal
componentr eht )h( ,ygrene nocomponent,
itaidar eht )(f)g( the,hcJ.timagneticp
tengameht (g))Thef( ,thetnSEDsenradiation
opmoofc ourpitch,
energy,(h) arethe ratioof the radiation
C.ciGladstone.
resultsnormalized
by the energy to the sum of the gas and magnetic
pmet noitaidar eht oenergies,t erutare(i)pmand
et satheg ehratio
t f30
o oof
itather ehgast dntemperaturea )i( ,seigrentoe the30radiation temperature. All values
30 are time-averaged over t = 6−7 s. The white
X-ray
Figure3, while
the SEDs
ULXs
FigureComparison
SEDs
tingam esohw srotceandv ytiblack
colev arrows
eht etacin
idpanelni )a(luminosity
lenaindicate
p ni swothe
rshownra velocity
kcalbindnvectors
a
(a)
whose
magnitude
exceedofthe
escape velocity.
The7.dashedline in ofpanel(b)obtainedis the from our radiation hydrodynamic
simulations
withshowsthosetheof ULXs.
Topflux
panel: the SED for the simulation modelb re ppu eht morf dephotosphere,
rusaem ssenkat
cihwhicht arelacitpo eopticalht hcihw
tby
a ,ertheehp
sX-rayotohp luminosity
thethicknessmeasuredfrom the upperboundaryis unity. The arrowin panel (g)radiative
normalized
shown
in Gladstone40
40
40
2 and viewing angle i = 10◦ –20◦ and the SED of NGC 1313
.rotcev
vector.
with
Ṁ
≈
1000L
/c
E
et al. (2009). The SEDs obtained by our numerical simulations
X-2. Middle panel: Ṁ ≈ 200LE /c2 (i = 0◦ –10◦ ) and IC 342 X-1. Bottom panel:
are similar
to those of ULXs. The
observed X-ray spectra50of
50
50Ṁ ≈ 200LE /c2 (i = 40◦ –50◦ ) and NGC 5204 X-1. Red or blue solid curves
ULXs with photon index Γ " 2 can be fitted well by the models
display the SEDs calculated in this work. Black points with error bars show
60
with 60i " 30◦ . On the other hand, 60for ULXs with photon index Γ
the XMM-Newton EPIC pn data corrected for absorption, which are shownin Gladstone et al. (2009) and are provided by J. C. Gladstone. Each SED is# 2, 70spectra with 20◦ " i " 50◦ 70are fitted well.
70
normalized by its isotropic X-ray luminosity.
For IC 342 X-1 and NGC 5204 X-1, the soft X-ray component(A color version of this figure is available in the online journal.)
80
80 1 keV of our results is weaker
80
below
than that of the ULXs.

0(data; Gladstone 2009)

r/RS

10

20

30

r/R

40

50

0

10

10

20

20
60

60

70

70

2

50

50

2

S

R/z

R/z

40

2

S

30

S

40

80

fi

R/z

30

z/RS

z/RS

80

10

20

Super-Edd disk & clumpy winds
Some ULXs exhibit the time
variations of X-ray luminosity,
implying the launching of clumpy
out ows.

Middleton+11

Launching of clumpy winds isalso reported by observationsof NLS1s or V404 Cyg.

Jin+17 see also Motta+17輻射加速円盤風の発生と分裂
Takeuchi, Ohsuga, Mineshige 2013

Jet

Super-Eddington disk+ Jet

accretion

ClumpyOut ow

fl

 

fl

fl

Time-dependent, Clumpy out ow
with wide angle

RT instability

Clumpy out
Kobayashi+18

w
tflo

Ou

The size is ~100Rs.
Out ow velocity is ~0.1c.
Rotation velocity is 30% of Vkep.

 

 

fl

Black Hole

Absorption lines
Sheet like structure
Out ow velocity ~ 0.1-0.2
Size (azimuthal direction) ~ 100Rs
Rotation velocity ~ 30% of Vkep

Absorption lines
Out ow velocity of~0.1-0.2c agrees with theobservations ofblueshifted absorption
lines.
Pinto+16,
see also Kosec+18

Time variation
Sheet like structure
Out ow velocity ~ 0.1-0.2
Size (azimuthal direction) ~ 100Rs
Rotation velocity ~ 30% of Vkep

 

 

c

c

 

 

 

Our result is consistent with theobservations of ULXs
(Middleton+11) and V404 Cyg(Motta+17) in the case of
MBH~10-100Msun.

 

 

fl

fl

fl

Time variation
Timescale of the luminosity
variation (100Rs/0.3Vkep) is

BHスピン

Initial torus

Utsumi, Ohsuga, et al. in presee also Mckinney et al. 2014,
Sadowski et al. 2014,
Takahashi 2016

a=0.7

a=-0.7out ow

disk

RadiationenergyDensity &stream lines

50

50

a=0

0

z/rg

BHスピンの効果の調査

Single-loop poloidal
magnetic elds
(plasma-β < 100)

BH

Preliminary

50

0
r/rg

50

50

0
r/rg

50

BHスピン
600
400
200
-0.5

0

0.5

spin parameter (a)Mass accretion rate decreases withan increase of spin parameter (ifthe initial conditions are the same).

in ow velocity

angular velocity·
MBH /(LEdd /c 2)

50

0
r/rg

Preliminary
ISCO radius

1.0

0

Spin parameter

ω/ωKep−u r
Spin parameter

1.0

0
0

10

R[rg]

p

 

 

fi

fl

fl

The difference of accretion rate is caused by the difference of the ISCO
radius. Within the ISCO, the angular velocity drastically decreases andthe in ow velocity increases.

fl50

Preliminary

BHスピン
Energy conversion ef cienc(around the rotation axis)·
η = L/ MBHc 2z[rg]Poynting Flux
a=0

200

a=0.7

100
-0.8

-1.0

0

0.8

0

1.0

Spin parameter (a)Energy conversion ef ciency (η)
increases with an increase of |a|.
MAD（強磁場円盤）ならPoynting
Fluxはもっと増えるかも？

0

50

0

50

R[rg]

Poynting ux is drastically enhanced
for the case of spinning BH. It is
probably BZ effect.

BHスピン
Radio Image (top) and Time variability of X-ray and Radio

Liska et al. 2019

J. Miller-Jones et al.
V404 Cygniの観測：Super-Eddington天体の

fi

y

fi

fl

ジェットの方向が変わった（？）↑これのGR-RMHD版
が必要！
今後の課題

未解決問題；計算法による相違
Stone et al. 92
Jiang et al. 14a
Jiang et al. 14b

VET
Full Transfer

Ohsuga, Takahashi 16

INAZUMA コード（Asahina+ 2000）

厳密
近似2

Y.-F. Jiang et al.

Sadowski+ 14,15

a characteristic photon trapping radius rtrap (Begelman
through global numericalMcKinney+simulations yield14
promising reTakahashi,
Ohsuga13 et al. 2012).
Gonzalez
et al. 07
1978), within whichphoton diffusion
time is longer
than
sults
(Kawashima
Takahashi,
Ohsuga+16
inflow time so that the photons are accreted towards the
The simulation we present in this paper is designed
black hole with the gas before they have time to leave
to address these questions.
We solve
the full radiative
Takahashi& Ohsuga17
the system. Therefore, the model predicts that the totaltransfer equation in the Newtonian limit without adoptluminosity emitted by the disk is close to LEdd and only
ing any FLD or M1 like approximations (Jiang et al.
2014a). Due to the large computational expense of
depends on the actual accretion rate logarithmically inthe super-Eddington regime (Shakura & Sunyaev 1973).
our calculations, we have only completed one simulation
Therefore, the radiative efficiencies tend to be muchwith enough resolution to reach inflow equilibrium for
Takahashi,a Ohsuga,
lower than standard thin disks.
small radial range. This limits our ability to makeOhsuga
et
al.
09,
11
Farris etbut
al.still
08allows us to
Since the magneto-rotational instability (MRI)et al.detailed,quantitative predictions,13(Balbus & Hawley
1991;
Hawley et al.
1995;
identify the physical mechanisms that govern the flow.
Our primary result is that magnetic buoyancy signifiBalbus & Hawley 1998) is now believed to be thephysical mechanism responsible for angular momentum
cantly increases the vertical energy transport in thesetransfer, studying the accretion disk properties withsuper-Eddington flows. Therefore, they achieve radiaself-consistent MRI is the next step to significantly imtive efficiencies nearly as high as standard thin disks
prove our understanding of super-Eddington accretion(Shakura & Sunyaev 1973). A radiation driven outflow
flows. Because photons control the dynamics in the raalong the rotation axis is observed, but we see no evidence of photon bubbles (c.f. Begelman 2002) and beamdiation pressure dominated flows, an accurate numericalalgorithm to solve the radiative transfer equation is also
ing of the emission is mild (c.f. King et al. 2001).
another essential ingredient. Properties of the saturation
2. EQUATIONS
state of MRI in both gas pressure and radiation pressure
dominated regimes have been studied extensively with
The ideal MHD equations coupled with the time
local shearing box simulations (Stone et al. 1996; Turnerdependent radiative transfer equations we solve are
2004; Hirose et al. 2006, 2009b,a; Jiang et al. 2013b).
(Jiang et al. 2014a)
The magnetic field is not only able to provide the stress∂ρ
we need to transfer angular momentum, but also gen+ ∇ · (ρv) = 0,
erates an additional cooling mechanism (Turner 2004;∂t
Hirose et al. 2009b,a; Blaes et al. 2011; Jiang et al.
∂(ρv)+ ∇ · (ρvv − BB + P∗ ) = −Sr (P ) − ρ∇φ,2013a) and coronae above the disk (Jiang et al. 2014b).
∂t
Structures of accretion disks in the super-Eddington∂E
regime havediagram
also been
two (2D)
Figure 2. Space-timeof studied
density with(top),global
gas temperature(middle)componentofr (E)
magnetic
field (bottom) at r =+ ∇ and· [(E azimuthal+ P ∗ )v − B(B· v)] = −cS− ρv · ∇φ,as well three (3D) dimensional magneto-hydrodynamic √ ∂t(left) and r =20r(right).
Units
for
ρ,
T
and
B
are
ρ
,
T
,
P
.
The
white
lines
at
the
top
panelsshow
theapproximate locatios000φ equation is∂B(MHD) simulations, where radiative transfer
electron scattering
photospheremeasuredsurfaces of the disk.
− ∇ × (v × B) = 0.
(1)solved with
flux-limited
diffusion from(FLD)theor nearby
M1 closure∂t(Ohsuga & Mineshige 2011; McKinney et al. 2014;
O(v/c)の輻射輸送方程式を直接解く
rium and willbe thefocus of our analysis. Figure 4Sa̧dowski et al. 2014). These global simulations confirm!
"∂I shows that startingar T 4
many properties of the slim disk model. Particularly,
from∼ 4r , there is a significant+ cn · ∇I = cσa− I + cσs (J − I)sstrong radiation driven outflows are formed in these sim∂t ward mass flux
4π along the radial and vertical direct!
"
ulations (Watarai & Fukue 1999; Ohsuga & Mineshige
ar T 4
At 20r+s3n, Ṁ2013).
· vσina = 3.01
− JṀsum , Ṁout = −1.72Ṁsum w
4π
However, many questions remain to be answered. BeṀz = −0.29Ṁsum.
cause of the approximations made in FLD and M1, they
+ n · v(σa + σs ) (I + 3J) − 2σs v · Hcannot accurately capture the angular distribution ofv·v
v · (v · K)
the photons near the photosphere (Jiang et al. 2014,− (σa − σs )J − (σa − σs ).
ccin preparation). Since we have developed a more ac(2)curate numerical algorithm to solve the time dependent
radiative transfer equations directly (Jiang et al. 2012;
Davis et al. 2012; Jiang et al. 2014a), we repeat these#$calculations without adopting previous approximations.
Sr (E) = σa ar T 4 − Er
Although the slim disk model correctly determines thatv
advection along the radial direction is more rapid than
+ (σa − σs ) 2 · [F r − (vEr + v · Pr )] ,cradiative diffusion along the vertical direction, advec(σs + σa )tion in the vertical direction is not considered. Nor have
Sr (P ) = −[F r − (vEr + v · Pr )]global numerical simulations identified vertical advectionc$v #as playing a significant role. In contrast, it has previously+ σa ar T 4 − Er .
(3)been speculated that vertical energy advection associc
ated with buoyant magnetic field may exceed transport
Here, ρ, B, v are density, magnetic field and flow velocby photon diffusion (Socrates & Davis 2006) and suchity, P∗ ≡ (P + B 2 /2)I (with I the unit tensor), P is thetransport has been demonstrated in local shearing box
gaspressure, and the magnetic permeability µ = 1. The
simulations (Blaes et al. 2011; Jiang et al. 2013a). It is
total gas energy density isalso notable that the standard
slim disk
have diffiJiang
etmodelsal. 2014
culty fitting the spectra of ULXs (Gladstone et al. 2009),
B2
1
,
(4)
although sophisticated calculations of radiation transfer
E = Eg + ρv 2 +2
2

M1

Eddington/FLD

Newtonianor O(v/c)

Specialrelativity

Generalrelativity

未解決問題；計算法による相違

FLD近似やM1法と大筋で同様の結果が得られた.

ただし, 定量的には結果が異なる. なぜ？

Figure 3. Snapshot of disk structures for density (left) and radiation energy density (right) at time 1.13 × 104 ts . Units for ρ and
Er are ρ0 and ar T04 respectively.

˙! r

未解決問題；計算法による相違
M1
INAZUMA

M1

INAZUMAINAZUMAだと，回転軸に光が到達しないという問題が生じない
より正確な輻射場はINAZUMAで！

未解決問題；円盤全域の構造計算領域の拡大
M1

過去の研究
Slim disk

Slim

標準円盤

BH

計算領域を拡大することで,ようやく降着円盤の全体像がわかってきた. アウトフローのパワーも先行研究と違う！

Kitaki et al. 2001

未解決問題；円盤全域の構造

Super-Eddington diskは３つの領域に分かれる
ことが判明

ここまでのまとめ• NLS1s, ULXs, micro-quasarsのメカニズムや, SMBHの急速成長に深く関わるのがSuper-Eddington disk

• Super-Eddington diskの研究は, Radiation-MHD（RMHD）
計算により大幅に発展した. さらに, GR-RMHD計算による研究が進められている.

• BZと輻射, どちらがジェットを加速するのかを解明する
必要あり！計算法による違いや, 円盤の全体像を解明す

 ることも当面の課題.

図 6-2

1/2 ページ = 左右 80 ミリ × 天地 70 ミリ

標準円盤

スリム円盤

ブラックホール
降着率

ライン吸収による輻
射力やガス圧による円盤風が発生エディントン限界

標準円盤

ライアフ

←ここ

Ultra Fast Out ows (UFOs)
absorption line

Tombesi et al. 2011

• Blueshifted Fe absorption lines aredetected in 40% of Sy galaxies• Typical velocity is 0.1-0.3c•Out ow rate ~ Accretion rat•Kinetic power ~ Jet power
Jet

e

fl

.

disk winds.

flOne of the importantcandidates of the originof the AGN FEEDBACK

Line-driven wind
UV radiation: Matter is accelerated by the line absorption(bound-bound transition) of heavy ions with low ionization state.
X-ray: Strong X-ray irradiation prevent the out ow because theLine driven wind
metals are fully ionized.

Line driven wind
Failed
wind

-ray

X
BH

FailedV windUayX-r
Accretion
disk
UV

R < a few
BH10Rs

~100RsR < a few 10Rs

>>Accretionseveral 100R
disk s
~100Rs>> several 100Rs

The wind is launched,Line force does not
The matter cannot besince X-ray iswork, since heavy ions
The windaccelerated
is launched,by the line
attenuatedbyLine force doesnot
The matter cannot be
are highly ionized bysince X-ray
force,isbecause UV
work, since obscuration
heavy ions via theaccelerated
by the line
the strong X-ray. attenuatedby is weak.
radiation
failedare highly ionizedbywind. force, because UV
obscuration via thethe strong X-ray. radiation is weak.
failed wind. Line-winds
Nomura et al. 2016, 2017, 2020(see also Proga et al. 00, 04)

From the standard disk,
the disk wind is launchedby the radiation force forspectral lines (line-force).

BH(108Msun)
Line windsmStandard disk

MBH~108Msu
Ldisk~0.5LEdd

 

103Rg

Simulated spectra
Mizumoto et al. 2021

Density, ionization parameter, velocity
drastically change, and there are ve
absorbing regions along the line of sight.

Complicated Fe absorption
lines (H-like and He-likeiron) would appear.

Simulated spectra

Absorption lines from H-likeand He-like iron are resolved
by XRISM.

More detailed absorption
pro les can be understood byAthena.

fi

fi

Mizumoto et al. 2021

Thermal-radiative wind
Resulting density pro leby radiation-HD simulations

several 100km/s

105Rg

Gas is heated up via theComptonization and blown awayby the gas pressure force and theradiation force.

Simulated absorption features (Fe XXVI Ly
α1 & Ly α2) are obtained by MONACO.

Tomaru et al. 2020

Thermal-radiative wind
Chandra/HETGS data with best- tting model Simulated spectrum of a 30 ks XRISM observation
Fe XXVI Ly α2

Simulation results are consistent
with observations of BH binary,
H1743−322.

Fe XXVI Ly α1

Separation of Fe absorption
lines due to velocity difference
would be detected by XRISM.

fi

 

fi

Tomaru et al. 2020

問題点
直接光のみ（散乱光などは無視）

真面目に輻射輸送

ProgaやNomuraが採用している方法*ただし流体は止めている

低電離領域

低電離領
Higginbottom et al. 2014

散乱光を考慮すると円盤風内部の電離度がかなり変わる！
散乱光入りの輻射流体計算が必要だが計算量が・・・・

ここまでのまとめ• 標準円盤から吹き出す円盤風は, AGNやXRBの吸収線の起原と考えられている.

• 輻射流体シミュレーションによる研究が進められていて, X線観測との比較による検証が予定されている.

• 原子物理と輻射輸送をより正確に扱う必要があ  るが, 当面は難しいかもしれない.

偏光X線観測
Cyg X-1の偏光度は8.6%以下（硬X線の偏光観測）

Chauvin et al. 2018

X線偏光輻射輸送計算コード開発中！

BHから銀河スケールへ
銀河/宇宙論スケール

ブラックホールスケールから
銀河/宇宙論スケールへ

ク
ッ
バ
ド
ー

荒田, 矢島+ 2019

AGN/銀河スケール

フィ

ブラックホールスケール(c)和田/NAOJ

(c)高橋

給
供
量

質

銀河/宇宙論スケールからブラックホールスケールへ

まとめ• 精密な円盤＋ジェット＋円盤風の構造解明のための高精度大規模シミュレーション• 理論と観測の協働によるアプローチ• 宇宙の進化過程の解明に向けたBHスケールと銀河・宇宙論スケールとの連結