1 / 62

Computational challenges in astrophysics: planets, massive black holes and dwarf galaxies

Computational challenges in astrophysics: planets, massive black holes and dwarf galaxies. Lucio Mayer University of Zurich and ETH Zurich. Collaborators Stelios Kazantzidis (KIPAC, Stanford University) Chiara Mastropietro (University of Munich) Piero Madau (Univ. of California Santa Cruz)

stacy
Download Presentation

Computational challenges in astrophysics: planets, massive black holes and dwarf galaxies

An Image/Link below is provided (as is) to download presentation Download Policy: Content on the Website is provided to you AS IS for your information and personal use and may not be sold / licensed / shared on other websites without getting consent from its author. Content is provided to you AS IS for your information and personal use only. Download presentation by click this link. While downloading, if for some reason you are not able to download a presentation, the publisher may have deleted the file from their server. During download, if you can't get a presentation, the file might be deleted by the publisher.

E N D

Presentation Transcript

  1. Computational challenges in astrophysics: planets, massive black holes and dwarf galaxies Lucio Mayer University of Zurich and ETH Zurich Collaborators Stelios Kazantzidis (KIPAC, Stanford University) Chiara Mastropietro (University of Munich) Piero Madau (Univ. of California Santa Cruz) Monica Colpi (Universita’ Milano-Bicocca) Victor Debattista (Univ. of Lancashire) Thomas Quinn (Univ. of Washington) James Wadsley (McMaster University) Simone Callegari (PhD student, Univ. of Zurich) Robert Feldmann (PhD student, ETH Zurich) Tristen Hayfied (PhD student, ETH Zurich)

  2. Outline Three examples of “expensive” numerical simulations • require the need of powerful parallel supercomputers • (millions of resolution elements, millions of timesteps) • set a new state-of-the-art in the modeling of astrophysical systems • allow to make a substantial step forward in our understanding of the origin of cosmic structures • Formation of giant planets - planet formation • Pairing of SMBHs in merging galaxies  galaxy evolution, accretion of SMBHs, gravitational waves (GR) • Evolution of dwarf galaxies in the Cold Dark Matter Universe  galaxy formation/evolution, cosmology

  3. Cosmology and Hydrodyamics with Conspirators: James Wadsley McMaster Univ. Joachim Stadel Univ. Zurich Tom Quinn Univ. Washington Lucio Mayer ETH Zurich/U.Zurich Fabio Governato Univ. of Washington Stelios Kazantzidis KIPAC Stanford Greg Stinson McMaster Univ. George Lake Univ. of Zurich Multi Platform, Massively Parallel treecode + SPH, multi stepping, cooling, UV background, Star Formation, SN feedback, radiative transfer. State-of-the-art calculations in cosmological structure formation, galaxy formation, planet formation, Solar System dynamics(for code description see Wadsley,Stadel & Quinn 2004).

  4. Distributed memory parallel supercomputers (MPI used) LeMieux (decommissioned 2007) 4096 processor AlphaEV6 at Pittsburgh Supercomputing Center (US) Now replaced by 4096 procs Cray XT3 “Big Ben” 500 procs Zbox2 (since 2005)+ 500 procs. Zbox3 (2007) at University of Zurich ~1800 procs. Cray XT3 + Cray XT4 at Swiss Supercomputing Center

  5. ~ 277 Extrasolar planets to explain…. Marcy & Butler 2006 Two theories for giant planet formation: Core Accretion (to stages)= coagulation of rocky planetesimals into 10 Earth masses solid cores + runaway gas accretion Disk Instability (one stage) = gas in protoplanetary disk collapses directly into gas giant planets. Solid core eventually forms later Planets around ~ 10% of surveyed (mostly FGK) stars, majority has minimum mass Mmin > Mneptune. For those with Doppler+transit detection we know they must be gas giants from their low densities (Guillot 2006)

  6. Giant planet formation via disk instability Fragmentation is rapid process - once Q ~ 1 Jupiter-size clumps form after a few disk orbital times/hundreds of years (Rdisk ~ 30 AU) (Boss 1998, 2001, 2002; Kuiper 1959; Cameron 1978). -- Massive self-gravitating protoplanetary disk (M ~ 0.1 Mo) When (1)Toomre parameter Q=vsW/pSG < 1and (2) gas cools fast, on a timescale comparable to the orbital timesmall density perturbations rapidly amplified by gravity and locally gravitational collapse is possible. Gravitational instability nonlinear process, need numerical simulation Initial Qmin < 1.5 Boss 2001 3D eulerian simulation Density map

  7. Are clumps formed by GI transient or long-lived structures? Can they collapse into proto-giant planets? (1)Need very high resolution to model gravity accurately at all scales (Rdisk ~ 30 AU while Rclump ~ 0.1 AU) and resolve huge density gradients + no restrictions on computational volume  spatially and temporally adaptive codes, i.e. SPH or adaptive mesh refinement (2)Need to model gas thermodynamics accurately. i.e. balance between heating and cooling processes because fragmentation depends on how fast gas can cool - effective equation of state or some implementation of radiative transfer

  8. Mayer, Quinn, Wadsley & Stadel (Science, 2002; ApJ, 2004): • Locally isothermal equation of state (=fast cooling, shorter than obital • time) until gas becomes dense and optically thick, then switch to adiabatic equation of state (=no cooling). • Showed that with resolution < 0.1 AU clumps become gravitationally bound protoplanets with masses 0.5-8 Mjupiter and orbital eccentricities comparable with typical of extrasolar planets (0.1-0.3) Disk grows slowly to a mass ~ 0.1 Mo accreting from the molecular cloud core for a few thousand years (Class 0 to Class 1 stage –see e.g. Yorke & Bodenheimer 1999) Each calculation used up to to 200K CPU hours on LeMieux at PSC and Zbox1/Zbox2

  9. Wengen code comparison (2005-2008) Mayer, Gawryszcak et al., 2008

  10. 3D SPH simulations with radiative transfer 60 AU • Fragmentation depends on how fast the disk cools -- crucial to model realistically radiative physics in the disk • Mayer et al. (2007) : Radiative transfer simulations with flux-limited diffusion with realistic (Rosseland mean) opacities (d’Alessio et al. 2001) + disk edge cools as blackbody with adjustable efficiency • Fragmentation still possible but less likely than with locally isothermal EOS or tcool ~ torb (see also Rafikov 2006). • Requires higher disk mass compared to runs with locally isothermal EOS or tcool ~ torb (~ 0.15 Mo rather than 0.1 Mo) • Movie: disk growing from 0.02 Mo to 0.15 Mo over ~ 50 Torb

  11. How does the disk midplane cool? Radiation or what? • Disk midplane cannot cool via radiation because vertical radiative diffusion timescale >> orbital time • Upwellings/downwelling with typical speeds ~ 0.1 Km/s (orbital velocities ~ 1 km/s at 10 AU) in overdense regions • Transport the heat from the midplane to the edge in ~ 30 years ~ Torb at ~ 10 AU -- then outcome depends on tcool/torb at edge • Turbulent cells are: • - intermittent • associated with superadiabatic entropy gradients, suggestive of • convection (Schwarzschild criterion for convection verified) 0.5 AU T=120 years

  12. Varying simulations parameters (Mayer et al., ApJL. 2007) RS=emitting area/geometric surface area • For conservative cases • (RS ~ 1-1.2) one needs • M > 0.15 Mo for • fragmentation to happen • Fragmentation sensitive • to: • mean molecular weight • efficiency of radiative cooling at optically thin boundary (~ tcool/torb) • Gas mean molecular weight • increases in strong spiral • shocks due to vaporization • of ice grains m=2.4, RS= 1.5 m=2.4, RS= 1.2 m=2.7, RS= 1 m=2.7, RS= 1.2

  13. Work in progress: once clump forms resolution increased further to follow its collapse. New resolution limit ~ 1 Jupiter radius (Rj). 1.5 Mj clump contracts to ~ 3 Rj in less than 100 years (tfreefall ~ 50 years), reaches Tc ~ 10000 K) 10 RJ 1 AU 1 AU 60 AU Density Temperature

  14. Merging of galaxies and supermassive black holes Preamble: for what we know so far, SMBHs (masses 105 – 109 Mo) do not live in random places, they live in the (spheroidal) nuclei of galaxies. Hence the formation and evolution of SMBHs must be tightly connected with galaxy formation and evolution

  15. Structure formation in a Cold Dark Matter Universe IS via mergers Standard LCDM Model (WMAP1) Wm ~ 0.27 Wl ~ 0.73 Wb ~ 0.044 Shown in is large N cluster simulation with the Zurich treecode PKDGRAV2 What happens to the SMBHs when galaxies merge? Do they also merge? To answer: couple gravitational dynamics of CDM halos with baryonic processes for formation + evolution of galaxies and SMBHs. Multi-scale problem! From 100 kpc (halo scale) to 1 kpc (galaxy scale) to < 1 pc (vicinity of the SMBHs) But in cosmological simulations resolution too low to model gas physics and star formation realistically at scales << 1 kpc, where the SMBHs reside

  16. Orbital evolution of SMBHs in galaxy mergers The dynamical evolution during the merger of two systems each containing a SMBH can be divided into four main phases (e.g Begelman, Blandford & Rees 1980;Merritt 2006): Phase 1:Black holes follow galaxy cores in which they are embedded. The cores (stars+dm+gas) sink by dynamical friction against the dark mater background. 100 kpc – 100 pc LISA (Laser Inteferometer Space Antenna) NASA/ESA Phase 2:After the cores merge a SMBH binary should form as the black holes lose angular momentum due to dynamical friction by stars+dark matter and/or gas in the merger remnant100 pc-1 pc Phase 3 - Slow hardening of the SMBH binary by three-body encounters with stars plunging along low angular momentum orbits or friction/torques against gas 1 pc -0.01 pc Phase 4 -Final coalescence of the two black holes when the binary separation becomes small enough (< 0.01 pc) for gravitational radiation to become important.

  17. Phase 2 and 3 ( 100 pc – 0.1 pc) – stellar background • N-Body simulations – follow mergers of two spherical stellar systems with a SMBH at their center (modeled as a very massive particle), only gravity included • Binary SMBH stalls soon after its formation, at scales ~1 pc, because • stars are ejected by the binary as it hardens suppressing dynamical • friction (Milosaljevic & Merritt 2002). • ----> last parsec problem • Caveat: possible loss cone refill (1) via scattering of stars towards the black holes by massive perturbers in the galaxy(e.g. molecular clouds – Perets & Alexander 2006) or • (2) owing to centrophylic orbits in triaxial potentials (Berczik et al. 2006)

  18. ….but observed galaxy mergers are dissipative, not collisionless

  19. Nuclear gaseous disks in merging galaxies Massive gaseous and stellar disks (M ~ 108-1010 Mo) with sizes ~ 100 pc found in the center of merging galaxies and merger remnants (Downes & Solomon 1998) Located where nuclear starburst and AGN emission occur ---> gas in the disks is probably fueling SMBHs AND the starburst (relative timescales unclear) Mrk231 Disk rotate + have random (“turbulent”) motions (with s > kT)

  20. Evolution of the gas component in major merger: formation of nuclear disk

  21. Multi-scale dissipative merger simulations ----- from ~100 kpc to ~1 pc Mayer et al. 2006, 2007 Splitting of SPH particles (Kitsionas & Withworth 2002, Bromm 2004): increase gas mass resolution by a factor of 8 and force resolution to 2 pc (softening of SMBH is 2 pc) 108 yr before merger completed > 1 million SPH particles within the central kiloparsec, res. 3000 Mo (>2 million CPU hours among several supercomputers for 3 years) 200 kpc 60 kpc Equal mass mergers with 10% gas

  22. Formation of nuclear disk in multi-scale simulation g=1.4 t0=5.12 Gyr Dt ~ 5x105 yr Mdisk ~ 3x 109 Mo Mayer et al. 2007, Science Boxes are 200 pc on a side

  23. Rapid binary SMBHs formation – in less than 106 years after the merger is completed, 50 times faster than with stellar background! -- gas drag governs SMBHs decay, not stellar drag 6 kpc scale 60 kpc scale Mayer et al. 2007, Science 160 pc scale

  24. Disk structure/SMBHs decay depends on thermodynamics • With Effective equation of state (EOS)P=r(g-1)U. g=1.3-1.4- appropriate for • nuclear gas densities in circumnc. starburst – based on Spaans & Silk (2000, 2005) Vs ~ 50 km/s, s ~ 100 km/s, vrot ~ 300 km/s Rotationally supported, turbulent disk, high density, SMBHs supersonic (VSMBHs ~ 200 km/s). Density, thickness, s compare well with observed nuclear disks (e.g. Downes & Solomon’98) Withg=5/3, adiabatic EOS. Equivalent to no radiative cooling (shut off by e.g AGN feedback)Vs ~ 200 km/s, s ~ 200 km/s, vrot ~ 150 km/s. Pressure supported cloud, low density, SMBHs nearly subsonic. face-on edge-on Boxes show the inner 200 pc 5 Myr after galaxies merge

  25. But last parsec problem not solved yet! New multi-scale simulations taken to 0.1 pc resolution ~ 6 pc box ~ 120 pc box Mayer, Kazantizidis & Escala, in preparation, burned already 106 CPU hours

  26. The origin of the darkest galaxies in the Universe Draco Fornax Carina, Mb= - 8 Fornax, Mb= -13 Darkest galaxies known = Nearby dwarf spheroidals (dSphs) • dark matter dominated (from their kinematics s2 >> GMstar/R) • faint, low surface brightness (Mb > -15, mB > 24 mag arcsec-2) • Low angular momentum content, v/s < 0.5 • Very low gas content: Mgas < 0.1 Mstar < 0.1 Mdark (Gallagher, Grebel & Harbeck 2003) MHI < 0.01 Mstar • Variety of SF histories, truncated or extended (Dolphin et al. 2006)

  27. The morphology-density relation in the Local Group • 2 Giant spirals • 40 Dwarfs 45% dIrrs • (Gas rich,vrot/s > 1 • Low surface brightness, • exp. stellar profiles) • 30% dSphs • 15% dEs • (Gas poor, vrot/s< 1, • Low surface brightness, • exp. stellar profiles) 10% transition

  28. TIDAL STIRRING of disky dwarfs In cosmological simulations w/hydro not enough resolution to study dwarfs ----> model interaction between a single dwarf galaxy and a massive spiral with hi-res N-Body + SPH sims, a few million particles per single dwarf model. Initial condition: a disky dwarf falling for the first time into the Milky Way halo (1) orbits and structure of galaxies/halos (NFW) from cosmological runs + scaling relations between baryonic disk and halo from Mo, Mao & White (1998) (2) free parameters (e.g. disk mass fraction, gas fraction in disk) chosen based on observations of late-type dwarfs (e.g. de Blok & McGaugh 1997; Geha et al. 2006) Mayer et al. (2000, 2001;2002)

  29. Tidal stirring = repeated tidal shocks at pericenters with primary galaxy (Weinberg 1994; Gnedin, Hernquist & Ostriker 1999)turnrotationally supportedlate-type dwarf (dIrrs) into faint spheroidals with very low v/s, < 0.5 (Mayer et al. 2001, 2002) Tidal heating/stripping of gas and stars + bar/buckling instabilities. Stellar bar tidally triggered, galaxy stable in isolation due to low surface density 1st orbit 2nd orbit 3rd orbit 4th orbit Tides do not strip enoughgas: in sims with SF final Mgas/Mstars ~ 0.1 (initial Mgas/Mstar ~ 0.3-0.5) while it is < 0.05 for dSphs (Mateo 1998). Complete gas removal requires additional mechanism…

  30. Tides + ram pressure in action Faber & Lin 1983; Mori & Burkert 2000: dSphs as ram pressure stripped dIrrs We include a tenuous gaseous halo in the MW consistent with observational constraints, <~ 10-4 atoms/cm3 and T ~ 106 K at 50 kpc(Sembach et al. 2003; Maller & Bullock 2004) Disky dwarf on typical “5:1” cosmological orbit: Apocenter = 250 kpc Pericenter= 50 kpc

  31. First time coupling cosmological simulation with individual interaction hydro simulation(Mayer et al., Nature, 2007) • Cosmological simulation of Milky Way formation: subhalos with masses and distances consistent with darkest dSphs (Draco, Ursa Minor) fell in at z > 1, when cosmic ultraviolet ionizing background (UV) 100 times higher than today (Haardt & Madau 1998, 2001) • Interaction simulation: dwarf galaxy with gas, stars and dark matter with same • orbit and mass as subhalos in cosmo simulation, include tides, ram pressure • and the UV background. • Dwarf model in • Interaction • simulation is • gas-dominated • because SF suppressed • by UV bg at z > 1 in low • mass halos Key result: Gas is completely stripped because strong UV radiation keeps it loosely bound to the dwarf

  32. Rapid removal of baryons produces very high M/L Dark matter and stars are only partially stripped -- suffer only tidal effects, no ram pressure+UV –- inner dark halo (within ~ 1 kpc) remains nearly intact. Final M/L > 100 naturally obtained starting from a “normal” mass-to-light ratio (~ 10) because all gas (=most of the baryons) is stripped by ram pressure+UV+tides

  33. If satellites fall in later, z < 1, some gas is retained and extended SF history results as UV bg fades T=8 Gyr =( gas/tdyn) (Kennicutt 1998) T=2.5 Gyr Leo I   From Hernandez et al. (2000) Star formation is periodic: gas driven bar inflow and tidal compression at pericenter Variety of SF histories in dSphs driven by different infall times

  34. What if the progenitor of some dSphs was gas dominated? NGC 2915 • Plausible assumption because: • Many late-type dwarfs (dIrrs) at z=0 have Mgas/Mbaryon > 0.5 today (e.g. McGaugh 2000; Geha et al. 2006; Mayer & Moore 2004) • (2)Both hydro simulations and analytical models of star formation naturally obtain that low mass disks should have a low star formation efficiency either because of feedback regulated star formation or because of inefficient molecular gas formation (Schaye 2004; Governato et al. 2007, see also talks by C.Wheeler and B.Robertson at this workshop).

  35. Orbital decay of SMBH binary in a spherical gas cloud • Driven by dynamical friction against the gaseous background + torque due to gas distribution around binary at scales < 1 pc • Evidence that binary SMBH continues to shrink below 1 pc Escala et al. 2004

  36. Comparison with spherical stellar cloud shows decay faster with gas if motion is supersonic and slower when it is subsonic (for the same density of the background). Ostriker 1999, Escala et al. 2004, 2005

  37. (1) SMBHs live in the nuclei of all massive galaxies (in their spheroidal component) (Kormendy 1999, 2003) (2) SMBH mass correlates weekly with luminosity of stellar spheroid and 1-D velocity dispersion of stellar spheroid (Magorrian et al. 1998) (Ferrarese & Merritt 2000, Gebhardt et al. 2000; Tremaine et al. 2002; Barth, Green & Ho 2004)

  38. On-going work (1): Recompute SMBHs orbital decay with realistic model for multiphase ISM in nuclear disk – beyond effective EOS Important because ISM model determines structural properties/thermodynamics of nuclear disk(also Wada & Norman 2001) Galactic-scale disk without (left) and with (right) blast-wave supernovae feedback(Mayer, Governato & Kaufmann 2008; Stinson et al. 2006)

  39. …and in rotating gaseous nuclear disks Simulations start from a model of a rotating gaseous nuclear disk ~ 100 pc in size (stellar bulge also included) with two SMBHs (MBH ~ 106-109 Mo) -- Non-thermal pressure support via effective equation of state (EOS) P = Kr g -- No real large scale mass distribution but resolution can reach below 1 pc (Escala et al. 2005; Dotti et al. 2006,2007 – see talk by Dotti) Main result: no stalling of SMBHs, decay continues down to resolution limit of the simulations due to dynamical friction (or ellipsoidal torque below 1 pc according to Escala et al. 2005) – no loss cone problem with gas Escala et al. (2005); Orbital decay with gas (black line) faster than with stars (green line) Dotti et al. 2006 (MBH = 4 x 106 Mo, res. 1 pc)

  40. Dotti, Colpi, Haardt, Mayer 2007 - see talk by Dotti Resolution reaches down to 0.1 pc with SPH particle splitting (Kitsionas & Withworth 2002, Bromm 2004; Kauffman, Mayer et al. 2006) • Decay continues but slows down • below 1 pc. • Extrapolating the decay rates found • by Dotti et al. (2007) and Escala • et al. (2005) at the resolution limit • (0.1 pc) one obtains that: • SMBHs with masses in the • range 4x106-5x107 Mo should reach • a separation of 0.01 pc in another • 106-107 yr • (2) Once at a~ 0.01 pc they would • coalesce via GW losses • in 107-1010 yr Last parsec problem” possibly solved? Does it make sense to extrapolate decay rate? Would gap formation be a problem?

  41. Nuclear disk simulations by Escala et al. (2004, 2005) and Dotti et al. (2006, 2007) have no direct link with the large scale dynamics and gasdynamics of the galaxy merger that produces the nuclear disk itself. need to make several assumptions for the structural properties of the nuclear disk (size, density profile, pressure support, kinematics etc..). Dynamical friction depends on such properties Next question: What is the timescale of binary SMBHs formation and of its subsequent sinking in a realistic dissipational merger? Strategy: multi-scale dissipational simulations of galaxy mergers, from 100 kpc to < 1 pc scales

  42. Conclusions • Prodigious gas inflows which result in the formation of nuclear disks at • ~100 pc scales are a natural consequence of dissipational galaxy mergers • Radial gas inflows driven by gravitational torques seen down to < 10 pc scales. These provide the gas reservoir to fuel SMBHs at sub-pc scales • The nuclear gas component in galaxies important for the merging of SMBHs at both large and small scales • at large scales (> 1 kpc) it can determine whether or not the cores of two galaxies will merge and the SMBH pair can form (in minor mergers) • -at small scales (< 1 kpc) it dominates the sinking of the pair and binary formation since SMBHs are embedded in massive, gaseous nuclear disks • gas can drive fast binary SMBHs formation in the nuclear disk (tbf ~ 106 yr in merger with only 10% gas) ~ 50 times more efficient than drag by stars • Both multi-scale mergers and nuclear disk simulations support fast decay driven by gas which may solve the last parsec problem

  43. Central disk structure and kinematics Disks exhibit significant non-axisymmetry (e.g. spiral arms, clumpy structures). Significant radial inflows down to ~ 1 pc scales. Inward radial velocities of 30-50 km/s sustained down to sub-pc scales the inner disk, > 107 Mo, could accrete onto the SMBH in less than 106 yr

  44. DOUBLE AGN EMBEDDED IN A STAR BURST NGC 6240 ULIRG DOUBLE BLACK HOLES VLBA RADIO GALAXY 0402+39 TWO COMPACT, RADIO FLAT, VARIABLE, DOUBLE NUCLEI 7.3 pc 100 million solar masses Claims of QSOs pairs with 3-10” separations Komossa review 2003 1.4 kpc

  45. MBH-Relationin Binary Galaxy Mergers Collisionless mergers indicate that when galaxies are negligible gas content (likely for low z) they move away from the relation (see also Ciotti et al. 2003). The tightness of the relation is consistent with the steep decline of merger rate below z = 1 (i.e. when galaxies are more likely to be gas poor) In dissipational mergers with star formation, the natural scaling between the amount of cold nuclear gas and the increase in s forces galaxies to move almost parallel to the relation explaining why they stay close to it. s: aperture dispersion that is measured within the effective radius Re averaging over different PA, inclinations and decompositions of the stellar mass distribution (e.g., Gebhardt et al. 2000). MBH: -In collisionless mergers we sum up the masses of the SMBHs if they form a pair or just the mass of the SMBH sitting closer to the center if the pairing process was aborted Assumption: Any close pair will actually merge (supported by Escala et al. 2004). Gas accretion based on both adiabatic and cooling simulations has to be viewed as upper limit. -In gasdynamical mergers we also add the gas mass around the SMBH pair within the resolution (~100 pc). Assumption: All the gas within the resolution will be accreted.

  46. Nuclear disks undergo intense star formation (obs, and large scale sims) Orbital decay time roughly independent on the fraction of gas/stars What matters is that the decay happens in a very dense nuclear disk produced by a gas-rich merger --- decay timescale naturally short. However once the binary is hard, the decay will continue or slow down depending on the role of gas vs stars below ~ 1 pc

  47. Semi-analytical models of co-evolution of galaxies and SMBHs -follow the merging history of dark matter halos and use simple prescriptions to decide which type of galaxy forms in halos based on angular momentum, density, star formation, type of merger -start from seed BHs of 100-200 Mo at z ~ 30 (collapse of Pop III stars) -put seed only in most massive halos at initial epoch -seed BHs grow during major mergers (mass ratio between galaxies > 1:3), due to coalescence with other BHs and spherical gas accretion (Eddington limited) - Black holes always merge when the galaxies merge (only evaluate phase 1!)

  48. Results of semi-analytical models of galaxy formation Evolution of optical luminosity function of AGNs M-sigma relation

  49. Powerful gas inflows at scales below 100 pc g=1.3 e=40 pc e=10 pc e=2 pc Gas inflow stronger with increasing force resolution because gravitational torques better resolved. Convergence approached at ~ 10 pc resolution.

  50. (4) Gravitational wave detection • Coalesecing compact objects are a primary source of gravitational waves. If two black holes with masses 106-109 Mo merge they will produce the strongest gravitational wave signals in the Universe. • The LASER INTERFEROMETER SPACE ANTENNA (LISA) will detect the merger of two SMBHs (104 – 106) Mo at z=1 with S/N of thousands. • For last year of inpiral detect masses, spin-orbit parameter, position on the sky very accurately (Cutler & Thorne 2002) • If SMBHs always merge when galaxies merge, ~100 events/yr at z~1, 1000/yr for z > 5, similar to BH-BH merger rate in the Galaxy (seen by LIGO) (Phinney 2004, Sesana et al. 2006)

More Related