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THE FIRST SUPERMASSIVE BLACK HOLES?. Mitch Begelman JILA, University of Colorado. COLLABORATORS. Marta Volonteri (Cambridge/Michigan) Martin Rees (Cambridge) Elena Rossi (JILA) Phil Armitage (JILA). NEED TO EXPLAIN:. Ubiquity of BHs in present-day galaxies QSOs with M>10 9 M at z>6
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THE FIRST SUPERMASSIVE BLACK HOLES? Mitch Begelman JILA, University of Colorado
COLLABORATORS • Marta Volonteri (Cambridge/Michigan) • Martin Rees (Cambridge) • Elena Rossi (JILA) • Phil Armitage (JILA)
NEED TO EXPLAIN: • Ubiquity of BHs in present-day galaxies • QSOs with M>109M at z>6 • Age of Universe < 20 tSalpeter Eddington-limited accretion would have to: • Start early • Be nearly continuous • Start with MBH >10 – 100 M
The Rees Flow Chart Begelman & Rees, MNRAS 1978
One more time, with calligraphy… Rees, Physica Scripta, 1978
18 years later…with 4-color printing! Begelman & Rees, “Gravity’s Fatal Attraction” 1996
ORIGIN OF SEEDS: 2 APPROACHES • Pop III remnants • ~100 (?) M BHs form at z~20 • Sink to center of merged haloes • Accrete gas/merge with other seeds • Direct collapse • Initial BH mass = ? • How is fragmentation avoided? • Growth mainly by accretion of gas • Problems faced by both: • angular momentum transport • Can be exceeded?
WHAT ANGULAR MOMENTUM PROBLEM? • WHOSE LIMIT? EDDINGTON • SEEDS FROM DIRECT COLLAPSE? • OBSERVING SEED BH FORMATION
Gas collapses if DM gas gas Halo with slight rotation DM Dynamical loss of angular momentum through nested global gravitational instabilities “BARS WITHIN BARS” (Shlosman, Frank & Begelman 89)
WHEN IS GAS UNSTABLE? • Init. collapse with conserved ang. mom. • Mestel, rigid, exponential disk • Include NFW halo potential • Vary fraction of gas collapsing, fd • Choose halo ang. mom. parameter λspin from lognormal distribution • Stability threshold: f T/W > 0.25 • Christodoulou et al. 1995, f = form factor
rigid Mestel exponential Probability of instability, lognormal distibution Max. halo spin parameter for instability, as function of gas infall fraction STABLE UNSTABLE
Global ang. mom. transport(“Bars within Bars”): M yr-1 • Mass distribution isothermal, • Accumulated mass dominates for • BWB fails at ?
Local ang. mom. transport( ): M yr-1 viscosity parameter (~1: Gammie 2001) • If T const. or incr. inwards, likely to reach center FRAGMENTATION UNLIKELY TO STOP INFALL
TEMP./METALLICITY DEPENDENCE • Metal-free, H2-dissociated gas falls in • (Metal rich and/or H2 fragments?) • Inflow rate, mass accumulation HIGH • Infall requires H2 or pollution by Pop III • Inflow rate, mass accumulation LOW FAVORS DIRECT BH FORMATION FAVORS POP III STAR FORMATION
WHOSE LIMIT? SUPPOSE A SEED BH SETTLES IN THE MIDDLE OF THE ACCUMULATED GAS ACCUMULATED GAS Max. BH accretion rate is for the mass of the ENVELOPE BH
RAPID GROWTH OF A PREGALACTIC BH + FORMATION OF THE SEED? “QUASISTAR” ACCUMULATED GAS Pregalactic halo Could seed BH grow from ~10 to >105 MSol at ? (Begelman, Volonteri & Rees 06)
YOU’VE GOT TO BE KIDDING WHAT IS A “QUASISTAR”? • Self-gravitating structure laid down by infall • Radiation-dominated, rotation crucial • Pre-BH: • Entropy small near center, increases with r • Very different from the supermassive stars postulated by Hoyle and Fowler
THAT’S BETTER WHAT IS A “QUASISTAR”? • Self-gravitating structure laid down by infall • Radiation-dominated, rotation crucial • Pre-BH: • Entropy small near center, increases with r • Very different from the supermassive stars postulated by Hoyle and Fowler • Post-BH: • “Nuclear” energy source is BH accretion • Expands and becomes fully convective • Like radiation-dominated, metal-free red giant
QUASISTAR STRUCTURE : PRE-BH • Mass m* (M) increases with time M yr-1 • Core with • Envelope • Entropy increases outward – convectively stable • Rotation increases binding energy • Outer radius constant • Core radius shrinks • Nuclear burning inadequate to unbind star • Core mass ~ 10 M constant • When core temp. rapid cooling by thermal neutrinos
CORE COLLAPSE AND FORMATION OF ~10-20 M SEED BH + SUBSEQUENT ACCRETION AT EDDINGTON LIMIT FOR QUASISTAR MASS
QUASISTAR STRUCTURE : POST-BH • BH accretes adiabatically from quasistar interior • Adjusts so energy liberated • Radiation-supported convective envelope (w/rotation) • Central temp drops to ~ 106 K • Radius expands to ~ 100 AU • Convection becomes inefficient well inside photosphere, where , but photosphere temp. drops as BH grows • Teff < 5000 K opacity crisis
OPACITY CRISIS • Pop III opacity plummets below T~5000 K • Higher temp than for solar abundances
Mayer & Duschl 2005 Metal-free opacities
Mayer & Duschl 2005 Metal-free opacities Plausible range of photosphere densities Analogous to Hayashi track, but match to radiation- dominated convective envelope
Est. min. temp. Mayer & Duschl 2005 If Tphot drops below minimum (~5000 K), flux inside quasistar exceeds Eddington limit, dispersing it.
OPACITY CRISIS • Pop III opacity plummets below T~5000 K • Higher temp than for solar abundances • Convection is relatively inefficient • Transition to radiative envelope at T~55,000 K ~ 10 Tphot • κ ~ κes at transition, but much lower at photosphere • If Tphot < Tmin , flux exceeds Eddington limit at transition quasistar disperses • Connect BH accretion physics to quasistar structure, predict Tphot(mBH, m*)
CONNECTION TO BH ACCRETION Once limiting temperature is reached, … dispersal is inevitable (and accelerates)
QUASISTAR MODELS WITH “TOY” OPACITY Begelman, Rossi & Armitage 2007
QUASISTAR MODELS WITH “TOY” OPACITY Begelman, Rossi & Armitage 2007
GROWING THE BLACK HOLE • Eddington-limited growth: • Ignoring opacity crisis:M yr-1 • Super-Eddington growth to ~107Min 108 yr • Onset of opacity crisis: Teff ~ 5000 K when • Estimates sensitive to: • Value of Tmin • Details of BH accretion • Energy loss details: rotational funnel, photon bubbles….
DETECTING A QUASISTAR • Most time spent as ~5000 K blackbody • Radiates at Eddington limit for 105m5M • Max flux ~
JWST z=6 z=10 z=20 PEAK OF BLACKBODY
DETECTING A QUASISTAR • Most time spent as ~5000 K blackbody • Radiates at Eddington limit for 105m5M • Max flux ~ • Better to observe @ 3.5µm, on Wien tail • Corona/mass loss hard tail, easier detection
JWST z=6 z=10 z=20 WIEN TAIL – MASS LOSS MAY IMPROVE FURTHER
Nseeds (Mpc-3) 1 HOW COMMON ARE QUASISTARS? 100 per L* galaxy 1 per L* galaxy Cumulative comoving no. density of seeds …but their lifetimes are short
COMOVING DENSITY OF QUASISTARS No enrichment L* galaxies Lifetime = 106 yr Some enrichment Metal enrichment model from Scannapieco et al. 2003 Heavy enrichment
COMOVING DENSITY OF QUASISTARS All 104 K haloes 100 per L* galaxy Lifetime = 106 yr 104 K haloes with λ<0.02 1 per L* galaxy Metal enrichment model from Scannapieco et al. 2003
DENSITY ON SKY • 10-3 seeds Mpc-3 (1/L* gal.) 103.5 QS deg.-2 • JWST FOV ~ 4.8 arcmin2 ~ 1 per random field Could be ~10-100x more common • How to distinguish from other objects? • Colors: ~pure blackbody (not dust reddened) • Observe on Wien tail • No lines (distinguish from T dwarfs) • Unresolved (distinguish from nearby starbursts) • Clustering (like 104 K haloes)
QUASISTAR DENSITY ON SKY 10/JWST field 1/JWST field Lifetime = 106 yr
WHAT HAPPENS NEXT? • If super-Edd. phase extends beyond opacity crisis, BH seeds could be as massive as 106 M • Worst case: super-Edd. phase ends at ~103 M • 10 tSalpeter between z=10 and z=6 growth by (only) 20,000 BUT • Exceeding LEdd by factor 2 squares growth factor! • Mergers can account for factor 10-100 of growth
CONCLUSIONS I • Angular momentum transport is fast if global gravitational instabilities set in (“Bars within Bars”) • BH can grow at Eddington limit for the surrounding envelope, which can be cvcvcvcfor the BH • BH seed can form in situ from the infalling envelope itself (aided by νcooling) or can be captured Pop III remnant
CONCLUSIONS II • BH seeds grow inside a “quasistar” powered by BH accretion, with a radiation pressure-supported convective envelope • Min. Teff of quasistar is ~5000 K, lifetime is > 106 yr • Quasistars could be common and may be detectable by JWST