the final parsec: orbital decay of massive black holes in galactic stellar cusps a. sesana 1, f....
TRANSCRIPT
The Final Parsec: Orbital Decay of Massive Black Holes
in Galactic Stellar Cusps
A. Sesana1, F. Haardt1, P. Madau2
1 Universita` dell'Insubria, via Valleggio 11, 22100 Como, Italy
2 University of California, 1156 High Street, Santa Cruz, CA 95064
Como, 20 September 2005
OUTLINE
>Merging History of Massive Black Holes
>MBHBs Dynamics: the “Final Parsec Problem”
>Scattering Experiments: Model Description
>Results: Binary Decay in a Time-Evolvig Cuspy Background: the Study Case of the SIS
>Effects on the Stellar Population >Returning Stars
>Tidal Disruption Rates >Implication for SMBH Coalescence
>Summary
MERGING HISTORY OF SMBHs
Z=0
Z=20
(Volonteri, Haardt & Madau 2003)Galaxy formation proceeds as a
series of subsequent halo mergers
MBH assemby follow the galaxy
evolution starting from seed BHs
with mass ~100M⊙ forming
in minihalos at z~20
During mergers,
MBHBs will
inevitably form!!
SMBHs DYNAMICSSMBHs DYNAMICS
1. dynamical friction (Lacey & Cole 1993, Colpi et al. 2000)
● from the interaction between the DM halos to the formation of the BH binary● determined by the global distribution of matter● efficient only for major mergers against mass stripping
2. hardening of the binary (Quinlan 1996, Merritt 1999, Miloslavljevic &
Merritt 2001)● 3 bodies interactions between the binary and the surrounding stars ● the binding energy of the BHs is larger than the thermal energy of the stars● the SMBHs create a stellar density core ejecting the background stars
3. emission of gravitational waves (Peters 1964)
● takes over at subparsec scales ● leads the binary to coalescence
DESCRIPTION OF THE PROBLEM
We want MBHBs to coalesce after a major merger
Dynamical friction is efficient in driving the two
BHs to a separation of the order
The ratio can be written as
we need a physical mechanism able to shrink the binary
separation of about two orders of magnitude!
GW emission takes over at separation of the order
GRAVITATIONAL SLINGSHOT
Extraction of binary binding energy via three body interactions with stars
Scattering experiments (e.g. Mikkola & Valtonen 1992, Quinlan 1996)
N-body simulations
(e.g. Milosavljevic & Merritt 2001)
resolution problem
> More feasibles
> need a large amount of data for significative statistics
(eccentricity problem)
> warning: connection with real galaxies!
> initial conditions
> loss cone depletion
> contribution of returning stars
> presence of bound stellar cusps
SCATTERING EXPERIMENTS
Y
X
Z
> MBHB M1>M2 on a Keplerian orbit with
semimajor axis a and
eccentricity e
> incoming star with m* <<M2 and velocity v
>The initial condition is a point in a nine dimensional parameter space:
> q=M2/M1, e, m* /M2
> v, b, , , ,
Our choices:
> In the limit m*<<M2: results are indipendent on m
*
we set m* =10- 7M (M=M1+M2)
> we sampled six values of q: 1, 1/3, 1/9, 1/27, 1/81, 1/243
and seven values of e: 0.01, 0.15, 0.3, 0.45, 0.6, 0.75, 0.9 for each q
> we sampled 80 values of v in the range 3x10- 3(M2/M)1/2 < v/Vc < 3x102(M2/M)1/2
> we sampled b and the four angles in order to reproduce a
spherical distribution of incoming stars
> Tolerance is settled so that the energy conservation for each orbit is of the order 10- 2 E*
> Integration is stopped when:
> the star leave ri with positive total energy
> the integration needs more than 106 steps
> the physical integration time is >1010 yrs
> the star is tidally disrupted
We integrate the nine coupled second order, differential equations
using the explicit Runge-Kutta integrator DOPRI5 (Hairer & Wanner
2002)
> At the end of each run the program records:
> the position and velocity of each star
> the quantities B and C defined as:
C and B-C distributions vs. x, a rescaled impact parameter defined as
M2/M1=1 M2/M1=1
e=0 e=0
SEMIANALITICAL MODEL
We consider:
> a MBHB with a semimajor axis a and eccentricity e
> a spherically simmetric stellar background
> (r) = 0(r/r0)- is the power law density profile. (0 is the density at the reference distance r0 from the centre)
> f(v,) is the stellar velocity distribution.
is the 1- D velocity dispersion (in the following we will always consider a Maxwellian distribution)
C and B can be used to compute the MBHB evolution
Writing d2N(b,t)/dbdt=2 b(b,t)v/m* and (b,t)=
0 F(ba x,t) we find:
Weighting over a velocity distribution f(v,) we finally get
H is the HARDENING RATE
Similarly we find the equation for the eccentricity evolution
K is the ECCENTRICITY GROWTH RATE
Starting from the energy exchange during
a single scattering event we can write:
F(bax,t) is a function, to be determined, of the rescaled impact parameter x
and of the time t and depends on the density profile of the stellar distribution
Early studies (Mikkola & Valtonen 1992, Quinlan 1996) assumed F(bax,t) =1
i.e. they studied the hardening problem in a
flat core of density 0 constant in time!!
Warning: connection with real galaxies!
1- Almost all galaxies show cuspy density profiles in their inner regions
r - 0< <2.5
(n.b. faint early type galaxies show steeper cusps that giants ellipticals)
2- In real galaxies there is a finite supply of stars to the hardening process
LOSS CONE PROBLEM
1-HARDENING IN A CUSPY PROFILE
We consider a density profile
r -
where =- 1
> If >1, then
> The hardening rate is:
Hard binaries hardens at a constant rate
only in a flat stellar background!
Eccentricity Growth
K is typically small: eccentricity
evolution will be modest
2-MODELLING THE LOSS CONE CONTENT
Definition: the loss cone is the portion of the space E, J constituded by those
stars that are allowed to approach the MBHB as close as x a,
where is a constant (we choose = 5)
Given (r ) we can evaluate the mass in the
unperturbed loss cone as
and the interacting mass integrating
where
M2/M1=1
M2/M1=1
e=0
THE SINGULAR ISOTHERMAL SPHERE (SIS)
> we can factorize F(bax,t) F0 (bax) x (t)
> The umperturbed loss cone mass content is Mlc ~ 3/2 M 2
> We model, as a studing case, the stellar
distribution as a SIS with density profile
r is related to t simply as dr/dt=31/2
> The MBHB mass is chosen to satisfy
the M- relation (Tremaine et al. 2002)
1- MBHB Shrinking
2-Distribution of Scattered Stars
The loss of low angular
momentum stars
Partial loss cone depletion
~20% of the interacting stars
returns in the new loss cone
of the shrinked binary
Stellar distribution flattening
and corotation with the MBHB
Interacting star distribution
tends to flatten and corotate
with the MBHB
Ejected mass
The ejected mass is of the order
Mej ≈0.7M
3-The Role of Returning Stars
Total shrinking
The shrinking factor scales as (M2/M)1/2
and is weakly dependent on e
Total loss cone depletion
The inner density profile
flatten significatively
Final Velocity Distribution
4-Tidal Disruption Rates
A star is tidally disrupted if it approaches
one of the holes as close as the tidal
disruption radius rtd,i~(m* / Mi)1/3r
*
We can then derive the mean TD rate as:
N TD stars / hardening time
> The TD rate is extremely high during
the hardening phase (respect to TD
rates due to a single BH ~10- 4 star/yr)
> The high TD rate phase is
extremely short
Hard to detect a MBHB via TD
stars
5-Binary Coalescence
As the shrinking factor is proportional to (M1/M)1/2, writing af = x
ah, we finally get
e=0
e=0.9
e=0.6LISA binaries (104-107 M⊙) may need extra
help to coalesce within an Hubble time!!!
What can help ?
> MBHB random walk (e.g. Quinlan & Hernquist 1997, Chatterjee et al. 2003)
> Star diffusion in the loss cone via two body relaxation (Milosavljevic & Merritt 2001)
> Loss cone amplification (loss wedge) in axisimmetric and triaxial potentials (Yu 2002, Merritt & Poon 2004)
> Torques exerted on the MBHB by a gaseous disk (Armitage & Natarajan 2002, Escala et al. 2005, Dotti et al. in preparation)
M <105M⊙
Summary >We have studied the interaction MBHB-stars in detail using scattering experiments coupled with a semianalitical model for MBHB and steller background evolution including: >a cuspy time-evolving stellar background >the effect of returning stars
>H in the hard stage is proportional to a -/2
>K is typically positive, but the eccentricity evoution of the binary is modest
>Interacting stars typically corotate with the MBHB
>MBHB-star interactions flatten the stellar distribution
>A mass of the order of 0.7M is ejected from the bulge on nearly radial corotating orbits in the MBHB plane>LISA binaries may need the support of other mechanisms to reach coalescence within an Hubble time
Results
Future Prospects
Investigate the contribution of other mechanisms to the binary hardening
Evaluate the eventual role of bound stellar cusps
Include this treatment of MBHB dynamics in a merger tree model to give realistic estimations for the number counts of “LISA coalescences”