1. Properties of massive black hole mergers Marta Volonteri University of Michigan

2. What we want to know ‣ How and when BHs form ‣ How and when BHs accrete mass ‣ How and when BHs merge ‣ How fast BHs spin Dotti + Sesana’s talks Dotti’s alk

3. What we want to know ‣ How and when BHs form ‣ How and when BHs accrete mass ‣ How and when BHs merge ‣ How fast BHs spin

4. First black holes PopIII stars remnants (Madau & Rees 2001, MV, Haardt & Madau 2003) ✔ Simulations suggest that the first stars are massive M ∼ 100-600 M sun (e.g., Abel et al. Bromm et al.) ✔ Metal free dying stars with M>260M sun leave remnant BHs with M seed ≥100M sun (Fryer, Woosley & Heger) Runaway collapse: stellar dynamics (Devecchi & MV 2008) ✔ Form dense stellar cluster at low metallicity ✔ Stars merge into a “super-star”, collapse into BH Direct collapse: gas-dynamics (e.g. Haehnelt & Rees 1993, Eisenstein & Loeb 1995, Bromm & Loeb 2003, Koushiappas et al. 2004, Begelman, MV & Rees 2006) ✔ Transport angular momentum on the dynamical timescale ✔ Cocoon of dense gas within which BH forms M BH ∼ 10 3 -10 5 M sun M BH ∼ 100 M sun

5. Mass function of seed MBHs PopIII remnants (MV, Lodato & Natarajan 2007, Devecchi & MV 2008) Proto-cluster Direct collapse

6. seed. SMBHS are grown from seed pregalactic BHs. These seeds are incorporated into larger and larger halos, accreting gas and dynamically interacting after mergers. time local galaxy high-z protogalaxies local galaxy high-z protogalaxies MV, Haardt & Madau 2003

7. GWs from MBH binaries Direct collapsePopIII remnants with Cutler, Berti & the LISAPE taskforce

8. Journey to the M-sigma relation MBHs along the hierarchical build-up of a massive elliptical galaxy seeds (direct collapse) Mergers happen in the most massive systems where MBHs are on the M-sigma (MV & Natarajan 2009)

9. MBHs along the hierarchical build-up of a massive elliptical galaxy seeds (PopIII remnants) MBHs move onto the M- sigma from below Journey to the M-sigma relation

10. GWs from MBH binaries Direct collapse PopIII remnants

11. What we want to know ‣ How and when BHs form ‣ How and when BHs accrete mass ‣ How and when BHs merge ‣ How fast BHs spin

12. Mergers and accretion Katzantidis et al., Mayer et al., di Matteo et al., Hopkins et al. Do MBHs accrete before or after merging?

13. Merger driven accretion  “Standard” numerical Bondi-Hoyle accretion  No feedback courtesy of S. Callegari

14. Merger driven accretion - turbulent motions - extrapolation from the refined simulation courtesy of S. Callegari High resolution simulations: M-σ without feedback

15. MBHs in circumnuclear disks Accretion gas particles are accreted only if their total energy (kinetic+thermal+potential) is less than a fixed fraction ε of the (negative) gravitational energy ( ε >0.5, accretion possible only resolving the Bondi radius of the MBHs) Dotti et al. 2009

16. Modulation in the orbiting MBH accretion rate Primary BH: ≈ 0.4 Secondary BH: ≈ 0.3 2.5 Myr) > ≈ 0.45 Co-rotating MBH (γ=5/3; h=0.1 pc) Dotti,.., MV 2009

17. MBH accretion rate depending on the sign of the angular momentum Primary BH: ≈ 0.25 Secondary BH: ≈ 0.1 3 Myr) > ≈ 0.45 Counter-rotating MBH (γ=5/3; h=0.1 pc) Dotti,.., MV 2009

18. MBH mass accretion MBH accretion rate depends on the relative velocity between hole and gas flow MBH accretion rate depends on the thermodynamical properties of the gas

19. What we want to know ‣ How and when BHs form ‣ How and when BHs accrete mass ‣ How and when BHs merge ‣ How fast BHs spin

20. a1a1 a2a2 L L a1a1 a2a2 a1a1 a2a2 L Low kick velocities (~100 km s -1 ) High kick velocities (~1000 km s -1 ) Recoiling MBHs

21. Random distribution of spin moduli in-plane spins spin-orbit isotropy aligned spins anti- aligned

22. Spin evolution accretion rate Simulations ang. mom. accreting flow Dotti,MV et al. 2009 Perego,.., MV 2009 Bardeen-Peterson effect → BH- SS disc alignment {

23. Spin evolution: Bardeen-Peterson effect Perego,.., MV 2009 see also Martin et al 2007; Lodato & Pringle 2006, 2007

24. Spin evolution Secondary BH initially counterrotating Primary BH a=0.9 a=0.1 a L ϑ Dotti,MV et al. 2009 angular momentum flip @ 2 Myr

25. Spin evolution a L ϑ Secondary BH Primary BH angular momentum flip @ 2 Myr CRE=cold disc,retrograde orbit HPE=hot disc, prograde orbit

26. retrogradeh ot prograde cold Recoiling MBHs Random distribution of spin moduli in-plane spins spin-orbit isotropy aligned spins anti-aligned spins

27. ejection rate mass increases with  redshift  spin-orbit isotropy

28. Hey, where are all these merging MBHS?

29. GWs: demography of “black” MBHs as a function of cosmic epoch Measure: masses, spins A multiwavelength effort: EM radiation: demography of “active” MBHs as a function of cosmic epoch Measure: luminosity, jets, host properties, spins

30. Hunting for sub-parsec binaries SDSS J092712.65+294344.0 (Bogdanovic et al. 09, Dotti,.., MV et al. 09) SDSS J153636.22+044127.0 (Boroson & Lauer 09, Chornok et al. 09?)

31. SMBH binary QSOs in the SDSS MV, Miller & Dotti 2009 JMM’s challenge: You predict all these tens of mergers per year that LISA will detect. How come there are so few sub- parsec binary QSOs in the SDSS?!? SAME black hole merger rate used for the LISAPE taskforce event rate predictions

32. Most MBH binaries are expected to occur at ‣ higher redshift ‣ lower masses ‣ lower luminosity than sampled by the SDSS quasar catalog MV et al 2003; Sesana, MV & Haardt 2007 SMBH binary QSOs in the SDSS Text λ=log(L/L Edd ) Not LISA sources! But, PTA sources! Sesana, Vecchio & MV 2009

33. ➡ MBH formation: determines how many mergers we detect ➡ accretion & feedback: determine how big MBHs are at merger ➡ mergers + accretion: determine how fast MBHs spin at merger, and the magnitude of recoil velocity