Supermassive black holes (SMBH) and formation of galaxies

Supermassive black holes (SMBH) and formation of galaxies Françoise Combes (Observatoire de Paris) SF2A-2003: PCHE

Relation entre masses BH-bulbes (BBR) Mbh = 0.2% Mbulge Bleu: vitesses stellairesVert: vitesses du gaz Rouge: disques de masers H2O, OH.. (Magorrian et al 98, Gebhardt et al 02, Ferrarese &Merritt 01, Tremaine et al 02, Shields et al 02)

Determination de la relation BBR -- mvts propres des étoiles (GC, Genzel et al, Ghez et al) -- raies d’absorption stellaires -- raies d’émission du gaz (sensibles aux outflows, inflows..) -- « reverberation mapping »: délai entre les variations du QSO dans les diverses longueurs d’onde (X, optique) => taille de la BLR IMBH: dans les amas globulaires M15, G1 dans M31 (IMBH: intermediate mass black holes) --M15< 103 Mo, pourrait être des remnants (van der Marel 2003) --G1, BH= 2 104 Mo, sur la BBR (Gebhardt et al 2002, HST-STIS) La BBR ne s’étend pas encore en dessous de ~107 Mo Mais s’étend aux QSO (le contraire par Wandel 99 non confirmé)

Relation with circular velocity Baes et al 2003 Circular velocity Link with dark matter?

Present density of BH and corresponding formation radiation There exists a BH in each bulge/spheroid, with a proportionality factor MBH = 0.002 Mb (BBR) rBH = 1.1 106 (MBH/Mbul/0.002) (Wbul/0.002h-1) Mo/ Mpc3 Optical radiation ropt = 1.4 105 (fbe/0.01)-1 Mo/Mpc3 fb fraction of light in B Hard X-ray total Hard X-rays 140 kev/s/cm2/st (z=1.5) rX = 3.8 105 (fXe/0.01)-1 Mo/Mpc3 Infrared ( assuming AGN contribution of 30%) rIR = 7.5 105 (fIRe/0.01)-1 Mo/Mpc3  The accretion radiation does not get out in optical light

Shaver et al 1996 Densité des quasars radio (Parkes flat-spectrum) Les quasars optiques suivent la même courbe très similaire à l'histoire de la formation d'étoiles

Gauche: Phases d'activité d'un quasar: rapide croissance, au taux maximum d' Eddington, rayonnement peu efficace puis phase active de 4 107 ans Droite: taux deformation des trous noirs, selon leur masse et z Haehnelt & Rees 1993

Haehnelt & Rees 1993 BH mass spectrum, integrated until z=0.5 Observations of Seyf and field galaxies are superposed Assumed relation between MBH and Vhalo MBH/Mhalo µ (1+z)6 e (-V/Vcir)

Interpretation of the BBR 1- QSO and stars main formation epoch coincide Feedback due to QSO outflows? (Silk & Rees 1998) Wind pushes the gas at constant v, with v3 fgs2 ~fw LE Condition for v > Vescape  MBH > 8 108Mo g (s/500)5 From 107 Mo, all the gas could be swept up, But the phenomenon is assumed spherical, in reality, jets are collimated, gas clumpy, and the gas is compressed to form stars… In any case, MBH ~ s5 (Faber-Jackson) ~ Mbulge

Sinking of SSC or IMBH 2- Sinking of SSC (Super Star Clusters) in a dark halo (Fu et al 2003) MBH/Mbulge = 10-4 only (like M33) Good profile and cusp expected

Are the BH the first objects to form? The first baryonic structures are ~106 Mo, like a GMC Then either they fragment, or fall in the center of dark haloes to form a BH as a single entity? Fragmentation must be inefficient, to account for the high frequency SMBHs in galaxies now, and the relative frequency of 109 Mo AGN at high z Pop III stars: could they go up to 200Mo? Radiative losses negligible at zero metallicity, also due to pulsational instabilities Above 260Mo, collapse to a BH directly At 105 Mo, BH automatically ==> IMBH at z=10-20 Madau & Rees 2001, Schneider et al 2002)

BH growth For simple dimensional relations, we can infer Racc = 0.3 M6/v22 pc and dM/dt is the Bondi accretion rate: dM/dt = 4 p R2 v r = (10-4 Mo/yr) = l M62/v23r since dM/dt ~ M2, then the accretion time is ~ 1/M. tacc ~108 yr /M8v23/r for very low BH this takes much larger than the Hubble time. Therefore it requires a large seed, mergers of BH, or very large densities, like in MW, 107 Mo/pc3 Accretion-dominated growth, tg = tacc. Nice for Seyfert 1 For QSO, they reach the Eddington limit, Ledd ~ M, the L ~ dM/dt ~ M2 L/Ledd ~ M, the BH growth slows down when approching Ledd. tedd = M/(dM/dt)edd = 4.5 107 yr (0.1/e) equating tacc = tedd, this occurs for Mt = 2 108 Mo v23/r (e/0.1)

IMBH: do they exist? Some theories predict them Observational constraints: lensing, X-ray sources, galaxy centers, if the BBR extrapolate? Globular clusters (M15?, G1 in M31) AGN in dwarf galaxies: NGC 4395(Filippenko & Ho 2003) MBH = likely 104-105 Mo (Seyf 1, no bulge) Low-ionisation, Lbol/LE = 210-2- 2 10-3 problem of dwarfs: host nuclear star clusters of ~106 Mo solution: only in the Local Group, possible to separate In M33 < 103Mo, factor 10 below the BBR

Double BH in the Milky-Way? Hansen & Milosavljevic 2003: evidence of an IMBH in the MW Brightest stars orbiting within 0.1pc of the SMBH are young massive MS stars, in spite of an hostile SF environment. (other solutions: star mergers, or exotic objects, Ghez et al 03) Stars were formed outside the central pc, and then dragged in by an IMBH (103-104 Mo) Relaxation time for normal stars is too large For IMBH, Trel = 1-10 Myr System SMBH-IMBH+stars, similar to comets with Sun-Jupiter Planetary migration Stars are dragged even after the star cluster has been disrupted

Radiation drag to explain the BBR Umemura et al (2001), Kawakatu & Umemura (2001) 3. Radiation drag driven accretion by bulge stars extracts angular momentum from the ISM and the gas accretes The MBH/Mbulge is then an universal constant = conversion fusion H --> He efficiency =0.7 % with a fudge factor fBH =0.3-0.5. The efficiency falls as 1/t2, with t the optical depth of the gas. But SF works in a clumpy medium Today this mechanism is inefficient, since E-galaxies and bulges have no gas!

Hierarchical models of galaxy formation 4. Hierarchical models explain very well the BBR (Haehnelt & Kauffmann 2000) The scatter is due to: -- Mgas of the bulge progenitor depends on s, but not on the formation epoch of bulge, while the M* depends on both -- mergers move the galaxies on the MBH- s relation, even at the end when there is only BH mergers (not much gas left to grow the BH) -- Bondi accretion depends of Ve, scatter due to scatter in Ve (Wang et al 2000; growth in MBH2, then 1/MBH) Saturates at 106 Mo if no external trigger

Kauffmann & Haehnelt 00 Evolution of the QSO density in the LCDM for MB< -26 Evolution of gas density in the LCDM model

Gas fraction in galaxies: 75% at z=3 and becomes 10% at z=0 Strong decrease of cool gas in DLAs (Storrie-Lombardi & Wolfe 2000) The gas fraction in major mergers is higher in fainter spheroids that form at high z (might explain the steeper cusp?) Remnant BH at z=0 MBH/Mbul = 0.2% (1% of gas has to be fueled to the BH) Elliptical/spheroids forming recently have smaller BH Rapid decline of QSO from z=2 to z=0 --decrease of merging rate --decrease of gas in galaxies -- BH accrete gas more slowly (tacc) LAGN/Lgal higher at high z , typical duration of QSOs 107yr Results

Multiple BH in the hierarchical scenario Merging and accretion history for BH growth Kauffmann & Haehnelt (2002) Typically a seed BH of 106 Mo forms at 5<z <10 and then gas is accreted, About 30 BH are merged. Today big elliptical’s BH accrete only by merging with small BH, but in the past, gas accretion was dominant. The slingshot effect is not efficient enough to reduce the separation of the two BH. But gas accretion is much more efficient, if the gas mass > the BH binary mass.  Explains the observed rarity of multiple BHs

Mergers of SMBH Milosavljevic & Merritt 2001

OJ287, light curve 100yrsPietila 98 3C75, Owen et al 1985 Roos et al 1993 VLBI maps of 1928+738 jet oscillations due to the orbital motions of the BH, period 3.2 yr

Destruction of bars Bars concentrate mass towards the center However, as soon as 5% of the galaxy mass is concentrated in the nuclear region, the bar is destroyed (in 2-5 Gyr) Rate is proven by observations Computations of the torque from the red image, on the gas distribution (Ha) Action on the gas: sign of the torques, depending on the phase shift between gas and stellar potential Exemple NGC 7479 NGC 7479 Nuclear bars accelerate the destruction of bars

Accretion of external gas simulated ==> Bar Renewal (Bournaud & Combes 2002) A galaxy is in continuous evolution, and accrete mass all along its life Several bar episodes can process in a Hubble time The Mbulge/Mdisk ratio and the gas fraction can evolve and therefore the morphological type can oscillate Without With accretion

With accretion Bournaud & Combes 2002 Without

Radial distribution of the accreted gas

Bars provoke their own destruction, by driving gas towards the center. After 5 Gyr there should not be any bar left Why so many bars today? (more than 2/3 of galaxies) Even more in NIR images Sample of 163 galaxies (OSU, Eskridge et al 2002) We computed the Qb, Q2, by Fourier Transforms of the potential Ratio of tangential forces to radial forces (torques) T= Ft/Fr(r,q) or T2= Ft(m=2)/Fr(r,q) Qb = max r,q T(r,q) Q2 = max r,q T2(r,q) (Block et al 2002) also Whyte et al 2002, axis ratio Statistics on bar strength N Qb

Quantification of the accretion rateBlock, Bournaud, Combes, Puerari, Buta 2002 Observed Doubles the mass in 10 Gyr No accretion

Bars have to be renewed -- Mass must be accreted in a large amount in the disk self-gravity of the disk has to increase with respect to the central concentration -- this mass must be gas, with dissipation: limited velocity dispersion will make the disk more unstable Interaction between galaxies are not sufficient -- they sometimes trigger bars, but also destroy them (Gerin et al 90) -- not efficient to reform bars, since the disk is not replenished with respect to the bulge and central concentration

Origin of the gas To have bars, continuous accretion at a large rate is required Dwarf companions: not more than 10% of accretion (interaction between galaxies heat the disk, Toth & Ostriker 92) Massive interactions: develop the spheroids But disks have to be replenished Required: a source of continuous gas accretion from the filaments in the near environment of galaxies Cold mode accretion, channeled by filaments (Katz et al, 2002; BDM? Pfenniger & Combes 94)

Several bar episodes in a galaxy disks, with secondary bars Regulation mechanisms Gas accretion at each gravitational instability The BH grows in parallel to the bulge from Combes 2000

Cusps and cores in E-galaxies Dichotomy: -- Cusps (power-law profiles) are the most frequent, characteristic of low-mass E-gal, with disky isophotes, some rotation --cores (2 power-laws) are rare, high-mass galaxies, high sv, boxy isophotes, no rotation Adiabatic growth of a black hole: gas accretion, destruction of nearby stars, etc..  produces a cusp (Cipollina & Bertin 94) Steep profile also from dissipative mergers, but core size not enough in dissipationless mergers Solution: mergers with heating due to binary BHs (Ebisuzaki et al 91)

Core size – MBH Ravindranath et al (2002) Milosavljevic & Merrit (2001) cusp slope g, versus m =MBH/Me solid line: Cipollina & Bertin 94 adiabatic growth Core galaxies with rb and ejected mass Mej = MBH (dash)

Remaining Problems The expected blue luminosity of AGN is too large: -- ADAF? (CDAF, ADIOS..) -- extinction, X-rays? how to lower the radiating efficiency? -- binaries: how to merge them more efficiently? -- cusp/core dichotomy -- why are the disks independent of the BBR? How is the BBR at high z? -- IMBH and threshold for the seeds -- M33-types: Mbh < 1000 Mo, 3 orders of magnitude below the BBR

Supermassive black holes (SMBH) and formation of galaxies