List Of Figures

List of Figures

	Figure 1: $20 cm$ VLA image of the radio source 3C 75 in the cluster of galaxies Abell 400. The image consists of two, twin-jet radio sources associated with each of two elliptical galaxies. The jets bend and appear to be interacting. The projected separation of the radio cores is about $7 kpc$ . Image courtesy of NRAO/AUI and F. N. Owen et al.
	Figure 2: Chandra X-ray image of the starburst galaxy NGC 6240, showing the two nuclear sources. Projected separation of the nuclei is about 1.4 kpc. Image courtesy of NASA/CXC/MPE/S. Komossa et al.
	Figure 3: A crude illustration of the parameter space for a SBH-IBH binary at the Galactic center. Assuming a circular orbit around a SBH of $6 3× 10 Mo .$ , a IBH with mass $MIBH$ and semi-major axis $a$ can be ruled out by measurement of an astrometric wobble of the radio image of Sgr A $*$ . The shaded regions show the detection thresholds for astrometric resolutions of $0.01$ , $0.1$ and $1$ milliarcseconds, respectively, assuming a monitoring period of $10$ years. The dashed lines indicate coalescence due to gravitational radiation in $6 10$ and $7 10$ years, respectively (From [84], see also [229]).
	Figure 4: Distribution of field star velocity changes for a set of scattering experiments in which the field star’s velocity at infinity relative to the binary was $vf0 = 0.5Vbin^ex$ . The binary’s mass ratio was 1:1, and the orientation of the binary’s orbital plane with respect to the $x$ -axis was varied randomly between the scattering experiments. Each plot represents $5× 104$ scattering experiments within some range of impact parameters $[p1,p2]$ in units of $a$ . (a) $[6,10]$ (b) $[2,4]$ (c) $[0.6,1]$ (d) $[0.4,0.6]$ . Solid lines in (a) and (b) are the distributions corresponding to scattering off a point-mass perturber. In (c) and (d), the mean of this distribution (which is very narrow) is indicated by the arrows. The gravitational slingshot is apparent in the rightward shift of the $dv$ values when $p$ is small, due to the randomization of ejection angles (from [133]).
	Figure 5: Mass ejected by a decaying binary, in units of $m12 = m1 + m2$ (solid lines) or $m2$ (dashed lines), calculated by an integration of Equation (24 ), with the coefficient $J(a)$ taken from [177]. Curves show mass that must be ejected in order for the binary to reach a separation where the emission of gravitational radiation causes coalescence on a time scale of $1010yr$ (lower), $109yr$ (middle) and $8 10 yr$ (upper).
	Figure 6: (a) Slices of the density $N (E, R,t)$ at one, arbitrary $E$ , recorded, from left to right, at $100$ , $101$ , $102$ , $103$ , and $104Myr$ (solid curve). Initially, $N (E,R, t) = 0$ for $R < Rlc$ and $N (E, R,t) = const$ for $R > Rlc$ . We also show the equilibrium solution of Equation (31 ) (dot-dashed curve). (b) The total number of stars consumed by the loss cone as a function of time (solid curve). The scale has been set to galaxy M32 with initial separation between the MBHs of $0.1pc$ . (From [152])
	Figure 7: Final density profiles from a set of 10:1 merger simulations in which each galaxy contained a black hole (a-d) and in which neither galaxy contained a black hole (e-g) [136]. The four thin curves in each frame correspond to four different pre-merger orbits. (a), (e) Space density of stars initially associated with the secondary galaxy; thick curves are the initial density profile. (b), (f) Space density of stars initially associated with the primary galaxy; thick curves are the initial density profile. (c), (g) Space density of all stars. Lower thick curves are the initial density profile of the primary galaxy, and upper thick curves are the superposition of the initial density profiles of the primary and secondary galaxies. Lines of logarithmic slope -1 and -2 are also shown. (d), (h) Logarithmic slope of the surface density profiles of the merger remnants. Thick curves correspond to the initial primary galaxy.
	Figure 8: Lagrangian radii around each of the two SBH particles in an equal-mass merger simulation [150]. From bottom to top, the radii enclose $-4 10$ , $-3.5 10$ , $-3 10$ , $- 2.5 10$ , $-2 10$ , $-1.5 10$ and $10- 1$ in units of the mass of one galaxy before the merger. The binary becomes “hard” at $t ~~ 11$ , and very rapidly heats the surrounding stellar fluid, lowering the local density.
	Figure 9: Evolution of the binary semi-major axis (a) and hardening rate (b) in a set of high accuracy $N$ -body simulations; the initial galaxy model was a low-central-density Plummer sphere [20]. Units are $G = Mgal = 1$ , $E = - 1/4$ , with $E$ the total energy. (a) Dashed lines are simulations with binary mass $M1 = M2 = 0.005$ and solid lines are for $M1 = M2 = 0.02$ , in units where the total galaxy mass is one. (b) Filled (open) circles are for $M = M = 0.005 1 2$ $(0.02)$ . Crosses indicate the hardening rate predicted by a simple model in which the supply of stars to the binary is limited by the rate at which they can be scattered into the binary’s influence sphere by gravitational encounters. The simulations with largest ( $M1, M2$ ) exhibit the nearly $N - 1$ dependence expected in the “empty loss cone” regime that is characteristic of real galaxies.
	Figure 10: Results from a set of $N$ -body integrations of a massive binary in a galaxy with a $r ~ r -0.5$ density cusp [205]. Each curve is the average of a set of integrations starting from different random realizations of the same initial conditions. (a) Evolution of the “mass deficit” (Equation 43 ), i.e. the mass in stars ejected by the binary. For a given value of binary separation $a$ , the mass deficit is nearly independent of particle number $N$ , implying that one can draw conclusions from observed mass deficits about the binary that produced them. (b) Evolution of binary eccentricity. The eccentricity evolution is strongly $N$ -dependent and tends to decrease with increasing $N$ , suggesting that the eccentricity evolution in real binaries would be modest.
	Figure 11: Observed surface brightness profile of NGC 3348. The dashed line is the best-fitting Sersic model to the large-radius data. Solid line is the fit of an alternative model, the “core-Sersic” model, which fits both the inner and outer data well. The mass deficit is illustrated by the area designated “depleted zone” and the corresponding mass is roughly $3 × 108Mo .$ [76].
	Figure 12: Effect on the nuclear density profile of SBH ejection. The initial galaxy model (black line) has a $r ~ r-1$ density cusp. (a) Impulsive removal of the SBH. Tick marks show the radius of the black hole’s sphere of influence $rinfl$ before ejection. A core forms with radius $~ 2rinfl$ . (b) Ejection at velocities less than escape velocity. The black hole has mass 0.3% that of the galaxy; the galaxy is initially spherical and the black hole’s orbit remains nearly radial as it decays via dynamical friction. The arrow in this panel marks $rinfl$ in the initial galaxy. [140].
	Figure 13: Final spin $^a$ of a remnant black hole in terms of its original spin, for mass ratios $q = 0.1$ (red), $0.3$ (green) and $0.5$ (blue). The change in spin was computed using the test-particle approximation for $L LSO$ [16, 91, 227]. Upper (lower) curves correspond to prograde (retrograde) capture from the equatorial plane; dashed curves are for capture over the pole. Capture of a low-mass secondary is likely to spin down the larger hole unless the latter is slowly rotating initially. Capture of a massive secondary results in spinup unless infall is nearly retrograde or the original spin is large.
	Figure 14: Steady-state spin distributions produced by successive capture from random directions at fixed mass ratio $q$ , for $q = (1/32, 1/16,1/8, 1/4,1/2)$ . Curves were generated using Monte-Carlo experiments based on the test-particle approximation, $q « 1$ ; hence the curve for $q = 1/2$ should be viewed as illustrative only.