super-critical accretion flow? shin mineshige (kyoto univ.) with k. ohsuga, k. vierdayanti, k....

Super-Critical Accretion Flow?Super-Critical Accretion Flow?Shin MineshigeShin Mineshige (Kyoto Univ.)(Kyoto Univ.)

with K.with K. Ohsuga, K.Ohsuga, K. Vierdayanti, K.Vierdayanti, K. Ebisawa, T.Ebisawa, T. KawaguchiKawaguchi

AGN conference @ Xian AGN conference @ Xian (10/16/2006)(10/16/2006)

1.1. General introductionGeneral introduction

2.2. Slim disk model (simplified model)Slim disk model (simplified model)

3.3. Observational tests - case of ULXsObservational tests - case of ULXs

4.4. 2D radiation-hydrodynamical (RHD) simulations2D radiation-hydrodynamical (RHD) simulations

1. Introduction 1. Introduction (background)(background)

It is widely believed that the luminosity It is widely believed that the luminosity ((LL) of any accreting objects cannot ) of any accreting objects cannot exceed the Eddington luminosity (exceed the Eddington luminosity (LLEE).).

・・What is the Eddington luminosity?What is the Eddington luminosity?

・・What do we know from observationsWhat do we know from observations ??

Spherical accretion system cannot shine at L > LE.

What is the Eddington What is the Eddington luminosity?luminosity?

radiationpressure

accretinggas

accretinggas

Gravity > rad. pressure

⇒ GM/r 2 >κF/c ⇒ L < LE=4πCGM/κ (∵ F=L/4πr

2)

Super-Eddington flux (F > LE/4πr2) is possible

because of radiation anisotropy (!?)

DiskDisk accretion may achieve L>Laccretion may achieve L>LEE

radiation

pressure

accretinggas

accretinggas

BH

Low-energy photons

trapped photons

Low-energy photons

High-energy photons

radiative diffusion & accretion

　 (c) K. Ohsuga

What is Photon trapping?What is Photon trapping? Begelman (1978), Ohsuga et al. (2002)

When photon diffusion time (Hτ/c) exceeds accretion time, photons are trapped within flow.

rtrap~ (Mc2/LE) rs .

UltraUltra LuminousLuminous X-rayX-ray sourcessources (ULXs)(ULXs) Makishima et al. (2000), van der Karel (2003)

Bright (~1040 erg s-1) compact X-ray sources Successively found in off-center regions of nearby galaxies. If L < LE, black hole mass should be > 100 Msun.

LE ~ 1038 (M/Msun) erg s-1

Two possibilities Sub-critical accretion onto intermediate-mass BHs (M>100Msun). Super-critical accretion onto stellar-mass BHs (M~3-30Msun).

IC342 galaxy

Narrow-lineNarrow-line SeyfertSeyfert 11 galaxiesgalaxies (NLS1s)(NLS1s)

What are NLS1s? Narrow “broad lines” (< 2000 km s-1) Seyfert 1 type X-ray features Extreme soft excess & variability

Seem to contain less massive BHs Good analogy with stellar-mass BHs in their

very high (large L ) state. High Tbb (∝MBH

-1/4) ⇒ large soft excess Small (GMBH/RBLR)1/2 ⇒ narrow line width

Boller et al. (NewA 44, 2000)

NLS1s = high L/LE

system

2. Slim disk model2. Slim disk model

Slim disk model was proposed for Slim disk model was proposed for describing high luminosity flows as an describing high luminosity flows as an extension of the standard disk.extension of the standard disk.

・・What is distinct from the standard disk What is distinct from the standard disk

model?model?

・・What are its observational signatures?What are its observational signatures?

Basics

This occurs within trapping radius

rtrap~ (Mc2/LE) rs

Model

Similar to the standard disk model; radially one-zone model but with radiation entropy advection

Slim disk modelSlim disk model Abramowicz et al. (1988); Watarai et al. (2000)

-2

-1

0

1

-1 0 1 2 3 4

log

L/L

E L=LE

log m≡log M/(LE/c2)

. .

accretion energy

trapped photons

.

3

8

4

9

)( 42gasrad T

dr

ssdTvr

partly

Slim-disk calculationsSlim-disk calculations Mineshige, Manmoto et al. (2002)

Low M rin～ 3rS ;Teff∝r

-3/4

High M rin～ rS

; Teff∝r -1/2

3 rS

.

M/(LE/c2)=1,10,102,103 　

　　　　　 MBH=105Msun

.

.

Slim-disk signatures 1.small innermost

radius2.flatter temp. profile

Wang & Zhou (1999), Watarai & Fukue (1999)

Disk spectra = multi-color blackbody spectra

Temperature profiles affect spectra

T ∝r -p ⇒ F ∝ν3-(2/p)

・standard disk (p =3/4)

⇒ F ∝ν1/3

・slim disk (p =1/2)

　　　　　　　 ⇒ F ∝ν-1

Spectral propertiesSpectral properties (e.g. Kato et al. 1998)

ν1/3

ν-1hν

Fν

hν

Fν

(small r)

(small r)

3. Observational tests: 3. Observational tests: Case of ULXsCase of ULXs

We examined the XMM/Newton data of We examined the XMM/Newton data of several ULXs which were suggested to several ULXs which were suggested to be candidates of intermediate-mass be candidates of intermediate-mass black holes.black holes.

・・ How to test the theory?How to test the theory?

・・What did we find?What did we find?

P-free disk modelP-free disk model (Mineshige et al. 1994)

Fitting with superposition of blackbody (Bν) spectra:

Fitting parameters:

Tin = temp.of innermost region (～ max. temp.)

rin = size of the region emitting with Bν(Tin)

p = temperature profile

⇒ Good fits to the Galactic BHCs with p=0.75

pr

r

rrTrTrdrrTBiFout

in

)/()( ;2)]([cos inin

Fitting with disk blackbody (p=0.75) + power-law

We fit XMM-Newton data of several ULXs ⇒ low Tin ～0.2 keV and photon index ofΓ=1.9　

If we set rin~ 3 rS, BH mass is MBH~ 300 Msun.

SpectralSpectral fittingfitting 1.1. ConventionalConventional modelmodel (Roberts et al. 2005)

NGC 5204 X-1

log hν

log conts

Spectral decomposition of DBB & PL components

DBB comp. is entirely dominated by PL comp.

Can we trust values derived from the minor component?

Problem with DBB+PLProblem with DBB+PL fittingfitting

log hν

log conts

NGC 5204 X-1

P-free model fitting, assuming T ∝ r -p

We fit the same ULX data but with p-free model only ⇒ high Tin ~ 2.5 keV and p =0.50±0.03 (no PL comp.)

MBH~ 12 Msun & L/LE~ 1, supporting slim disk model.

SpectralSpectral fittingfitting 2.2. p-freep-free modelmodel (Vierdayanti et al. 2006, PASJ in press)

NGC 5204 X-1

log hν

log conts

Why can both models give good Why can both models give good fits?fits?

Because the spectral shapes are similar below ~10 keV.

Both show fν ∝ν-1

in 0.1~10keV range.

power-law (Γ=2)

p-free with p=0.5

log kT (keV)

log Lx

Spectral fitting:Spectral fitting: SummarySummary (Vierdayanti et al. 2006, PASJ in press; astro-ph/0609017)

No evidence of IMBHs so far

P-free model fitting gives MBH <30Msun.

Low-temperature results should be re-examined!!

The same test is needed for NLS1s

4. Radiation Hydrodynamical 4. Radiation Hydrodynamical simulations simulations

The slim-disk model applies only to the flow with L ~ LThe slim-disk model applies only to the flow with L ~ LEE. For even higher L, we need radiation-hydrodynamica. For even higher L, we need radiation-hydrodynamical (RHD) simulations.l (RHD) simulations.

・・ What are the properties of the simulated flow?What are the properties of the simulated flow?

・・ What can we understand them?What can we understand them?

Our RHD simulationsOur RHD simulations Ohsuga, Mori, Nakamoto, S.M. (2005, ApJ 628, 368)

• First simulations of super-critical accretion flows in quasi-steady regimes. (cf. Eggum et al. 1987; Kley 1989; Okuda 2002; …)

• Matter with 0.45 Keplerian A.M. is continuously added through the outer boundary

→ disk-outflow structure

• Flux-limited diffusion adopted.

• Mass input rate: 1000 (LE/c2)

→ luminosity of ~3 LE

Injection

BH r/rs

z/r

s

500500

Initially empty disk

　 (c) K. Ohsuga

Overview of 2D super-critical Overview of 2D super-critical flowflow

BH r/rs

z/r

s

density contours & velocity fields

Ohsuga et al. (2005, ApJ 628,368)

gas energy radiation energy density density

disk flow

outflow

Case of M =10 Msun & M =1000 LE/c

2・

Why is accretion possible?Why is accretion possible?Ohsuga & S.M. (2006, submitted)

Radiation energy density is high; Erad ≫ EEdd≡LE/4πr 2c .

Then why does the radiation not prevent accretion? Note radiation energy flux is 　 　 Frad ∝(κρ)-1∇Erad.

→ High ρ and flat Erad profile 　　 result in weak Prad force.

Low ρ and strong Prad yield super-Eddington flux.

Gas & radiation flow:Gas & radiation flow: summarysummaryOhsuga & S.M. (2006, submitted)

In the disk region: large gas density flat Erad profile (+ photon trapping) → weak Prad force → slow accretion

In the outflow region: low gas & rad. density → strong Prad force → high-vel. outflow

The observed luminosity is sensitive to the viewing-angle; Maximum L ~ 12 LE !!

cos

luminosity

BH

Density contours

12

8

44D

2F(

)/L

E

our simulations

　 (c) K. Ohsuga

SignificantSignificant radiationradiation anisotropyanisotropy

⇒ mild beaming

Ohsuga et al. (2005, ApJ 628,368)

0

viewing angle

What are the causes of What are the causes of beaming?beaming?

Photon energy increases as θ decreases, why? - Because we see photons from deeper, hot region. - Because of (relativistic) Doppler effect.

Photon number increases as θ decreases, why? - Because of anisotropic gas distribution. - Because Iν/ν3 is Lorentz invariant & hν increases.

radiation

gasgas

Heinzeller, S.M. & Ohsuga (2006, MNRAS, astro-ph/0608263)

Further issues (1) Further issues (1) Line spectra of super-critical Line spectra of super-critical flowflow X-ray absorption in quasars? (Pounds et al. 2003)

• Outflow of 0.05-0.2c in X-ray spec.• Mass outflow rate ~ the Eddington rate• Can be explained by Prad driven wind. Same for BAL quasars?

Further issues (2) Further issues (2) Interaction with magnetic Interaction with magnetic fieldsfields

Magnetic fields are essential ingredients in disks• Photon-bubble instability (Begelman 2002) • Magnetic tower jet → global RHD+MHD simulations necessary

(Kato, S.M., & Shibata 2004)

ConclusionsConclusions

• Near- and super-critical accretion flows seem to occur in many systems (ULXs, NLS1s…?).

• Slim disk model predicts flatter temperature profile. Spectral fitting with p-free model proves the presence of supercritical accretion in ULXs. How about NLS1s?

• 2D RHD simulations of super-critical flow show super- Eddington luminosity, significant radiation anisotropy (beaming), photon trapping, high-speed outflow, etc.

L can be ~10 LE !!

• Issues: line spectra, magnetic fields, jet formation,…