AGN accretion and winds at parsec scales

Anton Dorodnitsyn	!
!

NASA/GSFC,	

!

Univ. of Maryland, IKI	Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

Collaborators:
Dorodnitsyn, A., KallmanT., Proga, D. 2014, 	

ApJL, (submitted)	

!

Dorodnitsyn, A., Kallman, T.	

ApJ, 2012, 761, 70	

!

Dorodnitsyn, A., Kallman, T., Bisnovaty-Kogan,G.S.,	

ApJ, 2012, 747, 8	

!

• Tim Kallman	

• Daniel Proga	

• G.S. Bisnovatyi-Kogan

Dorodnitsyn, Bisnovaty-Kogan,G.S., Kallman, T.,	ApJ, 2011, 741, 29	!

Dorodnitsyn, A., Kallman, T.	

2009, ApJ, 703, 1797	

!

Dorodnitsyn, A., Kallman, T.	

2010, ApJL, 711,112	

!

Dorodnitsyn, A., KallmanT., Proga, D. 2008, 	

ApJL, 657, 5	

!

Dorodnitsyn, A., Kallman, T., Proga, D.	

2008, ApJ, 687, 97	

!

Plan
AGN at 0.5-100 pc
scalesObscuration/Unification !
Models: disks/warps/winds/clumps ..!
Dusty Accretion and Outflows!
X,UV to IR: critical!
the rise of “Dynamical torus”!

AGN at sub-parsec
scales!
FLD numerical simulations!
of the torus illuminated by thequasar!
•Infrared-dominated flow!
•

Conclusions:!

What is the AGN unification?!

!

Active Galactic Nucleus

Seyfert 1, galaxy, NGC 5548

NGC 3277

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

AGN TAXONOMY I
Optical and UV lines

Optical
Case study: NGC 5548
• Narrow ~400km/s	

• Broad ~6000km/s

UV

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

Optical/UV

Seyfert II

UVProspects, challenges and evolution of AGN modeling in the ASTRO-H Era

AGN TAXONOMY I
Polarization[O III]

4969, 5007

[O III] 4363

H

H

NGC 1068

Miller, Goodrich & Mathews 91

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

AGN UNIFICATION• Seyfert 1: broad emission lines	• Seyfert 2: narrow lines only	So, “type 2” AGN seem to be “type 1” butwith no broad lines. Typically, Type 2 are lessluminous in optics
Astronomical Science

A significant fraction of AGN bolometric luminosity
comes in the 1 − 100µm infrared ”bump” in thedominatedby thespectral energy distribution (SED)(Elvisetouteral.component
1994; (Figure 1a), the absorption optical depth at
10 µm is moderate, !"0.4, whereas theRichards et al. 2006).
depth towards the inner component is almost five times larger at !"1.9 (Figure 1b,c, d). Obviously, most of the dust column
is located in the outer component.

••Interferometric mid -infrared observations of NGC
The dust torus in the Circinus galaxy
1068 (Seyfert 2), Circinus galaxy reveals multiThe Circinus galaxy at a distance ofcomponent structure
4 Mpc is the closest Seyfert galaxy. Itshows all signs of a classical Seyfert 2:
narrow allowed and forbidden emission lines, strong silicate absorption anda heavily absorbed X-ray spectrum. An
extended cone of emission line gas andthe presence of broad lines in the polarised optical flux (caused by scattering)
provide direct evidence that the central engine is hidden from our direct view
behind a substantial amount of dust.
Circinus is a spiral galaxy seen almostedge on. Thus, several magnitudesof visual extinction might be caused by
the dust lanes in the spiral disc, behind
which the nuclear region is located.

Warm (320K) - 2.1 parsec thick and 3.4 parsec indiameter. Inside of it - a hot >800K compact

Jaffe 2004

The high southern declination of Circinus(! = – 65˚) makes it an almost ideal target
for the VLTI: it can be observed for up to
12 hours during long winter nights, thus
allowing the projected baseline orientation between each of two UTs to swing
by up to 180˚ due to the earth’s rotation.
In five observing runs during the MIDI

200

100
DEC (mas)•

Meisenheimer K. et al., Mid-infrared Interferometry of Active Galactic Nuclei

0

–100

– 200
200

100

0
RA (mas)–100

Figure 3. The dust torus in the Circinus galaxy. The
left panel shows the smooth model (composedof two Gaussian brightness distributions, the rightpanel visualises the best-fitting patchy model.
Dashed lines indicate the opening angle of the ionisation cone.

200

100

0
RA (mas)

–100

– 200

component would also have to be inte
preted as a rotationally supported dis

The dust torus in the Circinus galaxy. !
Meisenheimer et al.	

ness with that of a black body leads to acovering factor of only 20 %. Moreover,
the observed correlated flux values arepoorly reproduced by the smooth Gaussian model, but rather seem to ‘wiggle’around it when plotting them as functionof baseline orientation. In order to test
whether a patchy brightness distribution
could improve the fit, we modified thesmooth Gaussian distribution by a foreground screen of randomly distributed
variations in transmission. A thousand
different screens were realised, the images were Fourier-transformed, and com-

It is worth noticing that the depth of thsilicate absorption towards the dust c
ponents in Circinus shows a different
haviour than that observed in NGC 10
In Circinus the depth of the silicate fea
ture towards the inner component is s
lower than towards the outer compon
Obviously the dust column through thouter dust component is not very highand the absorption trough is partly fillby silicate emission from the dust disc
The radio galaxy Centaurus A

The radio galaxy Centaurus A (= NGC
5128) plays a key role in extragalactic
astronomy: at a distance of only 3.8 M

Warm absorbers

4•The grating spectrographs on the X-ray telescopes Chandra and XMM-Newton provide unprecedented spectral
resolution up to ~ 10 keV.

!

• These show that X-ray spectra obtained from ~ half of low-red-shift active galactic nuclei (AGN) contain many

4lines from ions of Fe, Si, S, O, Mg, and Ne, and that these are generally broadened and blueshifted by 100-500
km/s
Kaspi, S., et al. 2002, ApJ, 574, 643; Steenbrugge, K.C. 2005, A&A, 432, 453!

•The presence of X-ray absorbing gas has been confirmed in the majority of AGNs which are bright enough toallow detections.

!

•There is also a partial correspondence between UV and X-ray absorbers
4

KINKHABWALA ET AL.

Vol. 575

FIG. 1 Illustrative 2-component global fit of an xstar model to the Chandra HETG spectrum
FIG. 1 ofIllustrative 2-component global fit of an xstar model to the Chandra HETG spectrum o
NGC 3783 (Ramirez et al, 2004). Vertical axis is counts/s/ Å, horizontal axis is observed wavelength
NGC 3783 (Ramirez et al, 2004). Vertical axis is counts/s/ Å, horizontal axis is observed wavelengt(Å). Solid (red) line is xstar model fit, (blue) pointers are positions of absorption lines in the(Å).model.
Solid (red) line is xstar model fit, (blue) pointers are positions of absorption lines in the modeg. 1.—Eﬀective area–corrected, first-order RGS 1 (red ) and RGS 2 (blue) spectra of NGC 1068 shifted to its rest frame (z ¼ 0:00379). The spectral disconies are due to chip gaps in the CCD arrays, bad pixels, and the previous inflight loss of one CCD for RGS 2 (" # 20–24 Å). Line labeling indicates the final
ion. All H-like (#) and He-like (r, i, and f ) principal order lines are labeled for each ion with the corresponding RRC edges indicated as well. Additionally,ance transitions (np!1s) are labeled as $ through % (short for Ly$–Ly% and He$–He%). Several Fe L-shell transitions are listed as well. Unlabeled feaat longer wavelengths (e.g., " ¼ 27:45, 27.92, 30.4, 31.0, 34.0–34.6, and 36.38 Å) are likely due to L-shell transitions in mid-Z elements. Line blueshifts
pecially noticeable at longer wavelength.

NGC 1068	

Kinkhabwala et al. 2002	

in 0>5 of each other, proving that possible spatial stratiion of diﬀerent emission zones is insignificant. The overpositional uncertainty between the PN and RGS
uments is less than the absolute pointing uncertainty of
elescope of 400 , implying a robust upper limit on the line
uncertainty of 8 mÅ (120 km s!1 at 20 Å).
gnificant velocity blueshifts from roughly 0 to 600 km

FIG. 1 Illustrative 2-component global fit of an xstar model to the Chandra HETG spectrum of

NGC 3783 (RamirezNGC
et al,
2004)	

3783 (Ramirez et al, 2004). Vertical axis is counts/s/ Å, horizontal axis is observed wavelength(Å). Solid (red) line is xstar model fit, (blue) pointers are positions of absorption lines in the model.

3.4. Radiative Recombination Continua
The spectrum also includes very distinctive radiative
recombination continua (RRCs) for H-like and He-like C,N, and O, which are produced when electrons recombine
directly to the ground state of these highly ionized species.
RRCs are broad features for hot, collisionally ionized

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

Hydrodynamical models
Simulated spectra
Rest Wavelength (Angstroms)
40

1.5

23.84

16.98

13.19

10.69

9.05

7.85

6.93

6.2

1.2

1.4

1.6

1.8

2

1

!=35°

0.5
1.5

!=47°

1

E FE

0.5

0.4

!=51°

0.2
0.1

!=55°
0.05
0

0.2

0.4

0.6

0.8

1

Rest Energy (keV)

Dorodnitsyn, A., Kallman, T.	

2009, ApJ, 703, 1797	

!

Dorodnitsyn, A., Kallman, T.	

2010, ApJL, 711,112	

!

Dorodnitsyn, A., KallmanT., Proga, D. 2008, 	

ApJL, 657, 5	

!

Dorodnitsyn, A., Kallman, T., Proga, D.	

2008, ApJ, 687, 97	

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

OBSCURATION• Quasi-static vs dynamical• Outflows/Accretion• Warps in accretion disk at a pc scale Phinney 89,

Sanders 89• Static IR-supported torus Krolik 07.
• Clumps/clouds: Elitzur 08. Krolik & Begelman 1988;

Beckert & Duschl 2004; Nenkova et al. 2008

• MHD winds from accretion disk, Konigl, Kartje 94,

Elitzur, Shlosman 06; Emmering et al. 1992; Everett
2005; Keating et al. 2012

X-ray polarimetry of warm
absorbers in Seyfert II

VS

Dorodnitsyn, A., Kallman, T. 2010, ApJL,
711,112Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

DYNAMICAL MODELS, NO RADIATION HYDROKeating et al. 12

The Astrophysical Journal, 749:32 (15pp), 2012 April 10

Wada

Keating

5

Figure 1. Schematic diagram of the structure of the radiatively accelerated
magnetocentrifugal wind model (Everett 2005). The sub-parsec-scale accretion
disk is the source of the continuum illuminating the wind. The wind has two
distinct components. The leftmost, thinner component is the “shielding gas,” andis modeled as a pure magnetocentrifugal wind. The wider streamline on the rightis the dusty wind, which is propelled by both magnetocentrifugal and radiative
acceleration. The black lines within the shield and streamline are examples ofthe zones in which Cloudy determines opacities and simulates photoionization.
In our calculations, the shielding gas is set to have an effectively negligible
column density so that the wind is dominated by the radiatively accelerated
component.

Figure 2. Gas streamlines for the fiducial model with NH, 0 = 1025
L/LEdd = 0.1, MBH = 108 M⊙ , αox = 1.6, and i = 60◦ . Streamlinplotted in the z-direction (along the axis of symmetry) in units of parsec
function of the r (radial) direction, also in units of parsecs. The strea
are shown for each iteration of the MHD radiative wind code. A
subsequent iteration (until convergence), the wind becomes more radradiative acceleration is taken into account.

Hopkins et al. 12

Dorodnitsyn et al. 08
Wada 2012sublimation radius, outside of which the silicate and graphite
dust can survive. The sublimation radius is dependent on a number of factors, including the bolometric luminosity and shape of
the incident continuum.
Schartmann et al. (2008) note during their investigations of
the dusty torus as a clumpy medium that it is important tohave a variety of sublimation radii for different sizes of dust
grains in order to accurately model the IR SED from the dust.
Their solution was to utilize a model with three different grain
species, each with five different grain sizes. Such a multicomponent grain distribution would be important, but given thecomplexity of our models and the large amount of time requiredfor computing each SED in even the simplified single-graincase, we will use a single grain type and sublimation radius in
our first approximation.
2.1. The Model CodeThe model code works in two separate steps. First, thestructure of the wind is determined semi-analytically, startingwith a pure MHD wind and then adding in radiation pressure
from dust and atomic lines. Second, a Monte Carlo simulation
uses the information about the structure of the wind to generate
an SED from 0.1 to 2000 µm. As we are interested primarily
in the more readily observable IR regime, we plot these resultsfrom 2 to 100 µm. Furthermore, at wavelengths < 2 µm, wecannot be certain of the accuracy of our SED calculations, as
the accuracy is affected by the parameters chosen for the Monte
Carlo simulation (see Section 3.1).
In more detail, the program starts by calculating the puremagnetocentrifugal wind dynamics (see Figure 1; Everett 2005)
without including the effects of radiation pressure. Parameterssuch as the black hole mass, column density at the base of thewind, and starting location of the wind are specified here. In thisgas along a streamline (this is described in detail in Ev
2005). Next, Cloudy (Ferland et al. 1998) is used to calcula
photoionization of the gas and determine opacities throug
the structure of the wind, given the central continuum. With
this information, the radiation pressure on the wind is calculThis is exactly as described in Everett (2005), except fo
addition of the opacity of dust, which is supplied via Clou
“ISM” dust model. This information is passed back into thestep, where the MHD wind structure is modified accordi
This process is iterated until convergence, which typicallyseven iterations.
Shown in Figure 2 are the gas streamlines for our fid
model (with parameters as specified in Section 4), plotted iz-direction along the axis of symmetry as a function of the r(r) direction. The streamlines are shown for each iteration o
MHD radiative wind code. The radiation pressure acts tothe streamlines radially outward, away from the central sowith each iteration, until the code converges.
The second stage of our calculation uses the Monte C
simulation program MC3D (Wolf 2003) to predict SEDs.
input for the simulation is the output from the MHD radi
wind code, which includes the structure of the wind and theof the central continuum source (we use the composite
from Richards et al. 2006, modified to remove IR emis
see Figure 3 and Section 3.3 for more details). This inp
passed to MC3D, which is a general purpose, three-dimens
Monte Carlo code that calculates heating and emission
given a central radiative source surrounded by a scattering
absorbing medium. MC3D first determines the temperature odust grains due to heating from the continuum and re-emi
from other dust grains. In a second step, MC3D’s ray trac
then used to create the SEDs from 0.1 to 2000 µm, whic
plot from 2 to 100 µm.

Temperature problem!
!

• One of the major problems which should be addressed by atheory of AGN obscuration is how the torus resists collapse intoa geometrically thin disk. 	

!

• If the torus is supported by rotation and gas pressure then thetemperature of such gas should be of the order of the virial
temperature: 	

!
Tvir,g=
2.6!

10 M6 /rpc K
5!

Clouds or clumps don’t help, as collisions drive T to Tvir, anyway	

cloud confinement problem is difficult.	

Clearly, such temperatures cannot be reconciled with theexistence of dust.

DUST SUBLIMATION RADIUS: THE OPACITY JUMP

Where is the inner edge?
Rsub ⇠ 0.6(Lbol /1046 )1/2 (Tdust /1400K) 2 pc.

!

•Eddington luminosity,
electron opacity

!
!
!

dust,

!
!

e

1044 M6

4
LUV
=
5
⇥
10
c,dust0.01 Ledd!
!

•Effective temperature of the
UV, XRAY to IR “conversion
layer”!
!

= 1.26!

•Very high opacity of UV on!

Ledd =

4⇥cGMBH

Teff =

4

GM⇥T ar2

⇥1/4

463 (

0.5 M7 1/4)2rpc

UV is absorbed in a
layer

!
!

•Critical IR luminosity in dusty
torus

l /R1pc = 1.3

Lc =

4⇥cGM
d

⇥ (0.030.1) Ledd

10 3 n7






!












 



POSSIBLE CONNECTION WITH BLR
Czerny & Hryniewicz: The origin of the BLR in AGN

beginning of the dusty torus, Rdust
et al. 2008). The ratio of the two raterms of the black hole mass, the aaccretion flow eﬃciency η (Lbol = η
1/6

r1000K
M8= 0.03
,
1/2
1/6
Rdust
ṁ η
0.1

Fig. 1.
The
BLR region covers the range of the disk with
Czerny etal.
14
an eﬀective temperature lower than 1000 K: the dusty windrises and then breaks down when exposed to the radiationfrom the central source. The dusty torus is the disk range
where the irradiation does not destroy the dust and thewind flows out.

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

where the black hole mass is expres
eﬃciency in units of 0.1. Since the dhole mass and the accretion rate isaccretion eﬃciency sources the outefactor of 30 than the inner radius.
tion rate drops by a factor of 10−3 ,the accretion eﬃciency drops by aour model the BLR is expected to r
(r1000K ≈ Rdust ). It is possible thatnation of the relatively narrow BLR
radio galaxies with clearly visible dline (see e.g. Lewis et al. 2010 and t
If indeed the scenario outlined incorrect interpretation, it means thatis not the dominant mode of outflow
other hand, this type of outflow is s
perhaps being responsible for finally
far away from the disk where high i

Acknowledgements. Part of this work wa
203 380136 of the Polish State Committee

Static radiation-supported torus
No. 1, 2007

AGN OBSCURING TORI

57

Fig. 2.— Solution with jin ¼ 0:5, " ¼ 1:5, ! ¼ 0:5, ## ¼ 10, and Q ¼ 3. Left: Radiation energy density. Right: Matter density. In both, the scale is logarithmic, and thewhite curve shows the surface on which #z ¼ 1, the photosphere on the top of the torus.

Krolik 07

Greater rotational support at the inner edge means that full
Whether this actually happens in practice is uncertain. In conKeplerian angular momentum is reached at a smaller radius. For
ventional global disk simulations, MHD turbulence very effiexample, if jin ¼ 0:75 but the other parameters are held fixed at
ciently redistributes angular momentum until it reaches a nearly
their fiducial values, rmax ¼ 1:76rin , and the matter density must
Keplerian radialprofile (De Villiers et al. 2003). If that happened
rise outward quite steeply (! ¼ "4:71) for there to be enoughhere, hydrostatic balance would be possible only if accompanied
optical depth across this diminished radial thickness; otherwise,
by a substantial positive : = F. This might perhaps be supplied by
the boundary condition matching the diffusion approximation enlocal sources of heat, such as stars or Compton heating due
 of the torus, but if sources
ergy density and the free-streaming energy density cannot be satto hard X-rays that penetrate the bulk
isfied. Conversely, if the fractional angular momentum support
such as these are inadequate, the disk would be forced to colat the inner edge is much smaller, rmax is much larger and thelapse to a thin configuration. On the other hand, in this context
density falls steeply outward: for jin ¼ 0:25, rmax ¼ 15:8rin andthe equilibrium angular momentum distribution
reached as a re
! ¼ 1:71.
sult of MHD turbulence might actually be that required by this
It is possible that distributed stellar mass may contribute sigmodel (eq. [9]) because the net ‘‘gravity’’ has been effectively
nificantly to the gravitational potential in the torus region. If so, reduced by the outward radiation force.
orbital speeds will decline more slowly with increasing radiusIn this context, it is notable that the velocity profile of the maser
than in the Keplerian prediction. When this is the case, the equaspots in NGC 1068
is /r"0.3, rather than /r"0.5as one might
expectfrom circular orbits in a point-mass potential (Greenhill
torial density profile able to produce an equilibrium switches from 
et al. 1996). Our model predicts a rotational speed that always
one that gradually falls outward to one that gradually rises. For
declines more slowly outward than simple circular orbits in theexample, for jin ¼ 0:5 and Q ¼ 3, ! drops from 0.5 to "0.19 as "Fromet roal. 2011
falls from 1.5 to 1.1.
gravitational potential would dictate;
for Dorodnitsyn
example, the mean
The optical depth parameter ## does not appear explicitly intational speed is constant as a function of radius when " ¼ 1:5.
the differential equations, so the shapes of the matter and energy
Thus, a shallow rotation curve may not signal a stellar contridensity contours, as well as the ! required by the radiation energy
bution to the gravitational potential — it could instead be a sympdensity boundary conditionat thephotosphere,
not dependofonAGN tomof radiation
Prospects,
challenges
anddoevolution
modeling
in thesupport.
ASTRO-H Era
it. However, the outline of the torus that results does change withMaser kinematics can also be used to estimate the central mass,iscous force simplifies to fvis, 'ppears in (3) reads

tr = 3/2µ⌦K . The viscous heating fu–8–

Effective
viscosity
fromgravito-turbulences is implied. Notice that in case of a geometrically thin equatorial part of the accretionthe viscous force simplifies to fvis, 'hat appears in (3) reads

tr

=
µ
t
@
v
.
v
ik
k
i
= 3/2µ⌦K . The viscous heating function,

umerical calculations we adopt the representation of (5) andal coordinates (e.g. Tassoul1978). The e↵ective (6)viscosity,v = µ tik @k vi .

ur numerical calculations we adopt the representation of (5) and (6) in2=L⌦,ndrical coordinates (e.g. Tassoul 1978). ⌫Thee↵ectivee↵crviscosity,⌫e↵ = L2cr ⌦,(7)n the critical spatial scale, Lcr which defines the size of the unstable pertur
nds on the critical spatial scale, Lcr which defines the size of the unstable perturbation.2
cancanbebe stabilized
byshear,thegivingshear,= G⌃/⌦is , or the perturba
cr perturbation
latterstabilized by the
Lcr = givingG⌃/⌦2 , orLthe
lized theby thee↵ectivee↵ective pressure,i.e. the velocity
dispersion,
givingdispersion,
the Jeans scale Lgiving
J =bypressure,
i.e. thevelocity
the Jeans sca
G⌃), where is the local dispersion velocity, Lin & Pringle (1987) argue that LJ < L <whereis thewelocal
velocity,
Linnumerical& Pringle(1987) argue that LJ
here, for simplicity
assumedispersion
that L ' LJ . Whenperforming
simulations
furtherin the text) we
L 'LLcr
not find any
noticeable
for simplicity
weexperimented
assume with
that
'and
LJdid
. Whenperforming
numerical simu
ence in the results.
er in the text) we experimented with L ' Lcr and did not find any noin the results.
2.2.

Photoionization equilibrium and X-ray heating and cooling
Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

EQUATIONS WE SOLVE

D=c ,where = 1/(⇢) is the photon mean free path, and  = /⇢. F(7)
one can see that in the di↵usion regime the radiation force does no
Viscous stress tensor with alpha viscosity/⇢. From
opacity.(??),
F (??), and (??)

D=c ,mean free path, and  =d depend on thegime the radiation
Dt + force· vdoes= 0,not explicitly
The cdi↵usion approximation tacitly assumes that optical deEnergytorus where IR pressure is importantRadiationindeedhas ⌧d > 1 where ⌧d i
Z I
1
1
1
E(r, t) =d I(r, , , t)cmixtureintheinfrared.
rp + ggas-dust
r
+
r
·
T,
rad
1

Dt v =
tacitly assumes that⇢optical depth ⌧ & 1. ⇢Most of the
If ⌧ optical< 1 the depthdi↵usionofapproximation
should be modified so a◆
ant indeed has ⌧✓
>
1
where
⌧
is
the
the
d e
d
2 that without such modification w
at
⌧
<
1
is
regained.
One
can
see
⇢Dt
= pr · v 4⇡ P B + c E E + n (⇤) + Qvisc ,⇢path, ! 1, and D ! 1, and |F| ! 1 which is in contradict– where7 – it should be |F| ! cE (??). That is, when optical dept
limit
✓
◆
mation should be modified so a correct limiting behavior
E⇢ ! 0, the standard di↵usion approximation is no longer applicab⇥D= ⇥ · Fwhen⇥v: P1 +
4 mean⇤P B freec⇤E E
hat withoutsuch
modification
⌧⌧
the
t
⇥
orderatofree-streaming
describe
correctly regions of small ⌧ , the standard
F| ! 1 which is in contradictionInwith
Radiation
Pressure
Z I
1RadiationFlux	

di↵usion
approximation(???). In this approximation
P(r, t) =d n̂n̂ I(r, , , t)
Z I
Fflux-limited=becomes
D rE,
(6)
c
(??). That is, F(r,when
opticaldepthsmall,orwhent) =–d 7n̂ I(r,– , , t)where ⇤F is the flux limiter. The flux limiter we adopted in Pap
oximationis isnolargelongerptical depth⌧ applicable.
1, theworkdi↵usioncoefficient
reads:
is that of
?:
0

1

1

0

0

regions ofFlux-limited
small ⌧ , thediffusion:
standard approach is to adopt the
2 + RLP⇤, (7)n (???). In this approximationis =replaced
c , by = ⇤F , ⇤F = 6 +(6)F = D rE, D
2
3RLP + RLP
limiterweadoptedin Paper
I,II
where
RLPand=|rE|/E.
One From
can see,(??),
that if(??),⌧ ! 0,
then
RLP !
=flux
mean
free
path,
andin
the
= current
/⇢.
and
(??)
is1/(⇢)large⌧is the1, photon
the di↵usion
coefficient
reads:

Levermore,& Pomraning 1981
⌧ radiation
1 RLP ! 0forceand ⇤doesF ! 1/3.
ee that in the di↵usion regime thenot explicitly
depend on theL=0.01Ledd

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

Density

Temperature

L=0.05Ledd

Prospects, challenges and evolution of AGN modeling in the ASTRO-H EraL=0.1Ledd
Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

4.7.

Mass-loss rate of the windAccretion and wind mass-loss rates
accretion from wind

The wind mass-loss rate, Ṁw is calculated at the upper, bottom and at the outer
Figure 8 (left) shows the time-dependent properties of this quantity.
– 27 – cylindrical boundary.
disk
accretion
If the entire torus is expelled within Nd dynamical times, with no material accreted, thenon average the wind mass loss rate satisfies hṀw i = Mtor /(Nd t0 ) ' 1.04 ⇥ 102 Nd 1 M yr 1 .
Mass-loss rates observed in the simulations are typically much smaller: In all runs the coreof the initial torus that contains most of mass survives complete evaporation (c.f. Figure 6).

Fig. 7.— Total accretion rate, Ṁa versus time, for di↵erent values L/Lddd . Notice, that the
Total withaccretion=rate,Ṁisa versus
time,modelfor different
values L/Lddd. a period of disk accretion.
model
0.01
the onlythat demonstrates
Afterdependenceof the total
wind
massFig. 8.— Time dependence of Timethe total
wind mass-loss
rate
, Ṁ
w for fiducial
Notice,
that
the
model
with
Γ
=
0.01
is
the
only
model
that
3
̇ w for fiducial Γ = L/Ledd (left)the end of this episode accretion stops for approximately
10 yrs.of the verticalloss rate, M(right).
(left) and3.9
the⇥energy
winddemonstrates a period of disk accretion. After the end of this
episode accretion stops for approximately 3.9 × 103yrs.	

= L/Ledd

and the energy of the vertical wind (right).	

From stops.
comparing
Figure 8 (left), the wind mass-loss rate mirrors theat t ' 2 ⇥ 104 yr the accretion rate declines and eventually
ThisFigure
is due7towith
truncation
accretionrate.
In the modelL=
0.01Leddandas the inner edge of an accretion disk progressesof the accretion disk at t ' 3.2 ⇥ 104 yr. That is, due
to rapidaccretion
the disk
depletes
towards the radiation source, the mass loading from the evaporation increases and produces
the reduced density cannot e↵ectively support screeningof its inner parts from overheating
an increase in Ṁw . Similar behavior can be noticed for other models.
and evaporation. After some time, hot gas fills the throat of the torus and accretion resumes
at much lower level through the capturing of the gas from the wind.

Accretion is interrupted for a period of 3.9 ⇥ 103 yrs, after which it resumes at the2

41

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

Torus covering factor and opening angle

Scatter plot showing the logarithm of the column density,
Ncol(cm−2). Each model is represented by a point. Colors
indicate different values of L/Ledd. Models shown from left toright: Γ = 0.01, 0.05, 0.1, 0.3 respectively.	

The logarithm of the number of models which at thegiven inclination have optical depth larger than one. Xaxis: the inclination angle in degrees measured from theaxis of rotation. Colors indicate different L/Ledd.	

Dorodnitsyn et al. 16(subm.)

Prospects, challenges and evolution of AGN modeling in the ASTRO-H Era

CONCLUSIONS!
•time-dependent radiation hydrodynamics calculations show,that for AGN luminosities greater than 0.01-0.1 Ledd external
illumination supports geometrically thick obscuration via
outflows driven by the radiation pressure of the reprocessed
IR radiation on dust.
!
• Radiation pressure on dust has a profound effect on AGN
pc-scale accretion and winds.
!
•The velocity of the wind can be comparable or greater thanthe escape velocity, or sometimes can be lower than theescape velocity suggesting that in some AGN a significant partof our torus is in the form of failed (at a larger distance) wind.
!