6/26/17
From the Dark Ages to the Cosmic Renaissance 1st BHI Conference
Harvard University High-redshift frontier Cambridge, May 2017 (t<1 Gyr)
Black Holes in the High-Redshift Universe
1st Stars st Volker Bromm 1 Galaxies Department of Astronomy University of Texas at Austin • First Stars (Pop III) à From Simplicity to Complexity
Structure Formation from Cosmological ICs Formation of a Population III Star: (Stacy, Greif & Bromm 2010, MNRAS, 403, 45) à Properties of First Stars, Galaxies, Black Holes
• ab initio structure formation (or: the Laplace Demon at work) Minihalo: 6 M ~ 10 Mo Input: Output: R ~ 100 pc Laws of nature Ω , Ω , Ω , B Λ m Stellar IMF, SFE, z ~ 20 σ8, ns, H0, Multiplicity… abundances, (supercomputer)
DM particle physics Black holes (Pre-stellar) Core:
3 M ~ 10 Mo ! Non-linear amplifier! R ~ 1 pc
1 6/26/17
How to make a star? Merging of Pop III Protostars (Greif, Bromm, et al. 2012, MNRAS, 424, 399)
• Gravitational instability (Jeans 1902) Gas Temperature • Gravity must overwhelm opposing thermal pressure: - AREPO (Springel 2010) Gas (Hydrogen/ - cosmological ICs Need: Helium) - completely ab initio: Mass > NO sink particles! Jeans mass (MJ)
Sir James Jeans
• Primordial clouds are hotter à need larger masses à first stars were more massive!
Population III : Radiative Feedback Pop III Star Formation: Growth by Accretion (Stacy, Greif & Bromm 2012, MNRAS, 422, 290) (Stacy, Greif & Bromm 2012, MNRAS, 422, 290) 16 A. Stacy, T. H. Greif and V. Bromm Stellar mass vs. time First Stars: Growth Under Feedback 13 age accretion rates over a longer period of time (50,000 yr) 1,500 yr 2,000 yr 3,000 yr and found that gravitational and Rayleigh-Taylor instabili- ties would continue to feed mass onto the disk and stars. We also note that, despite their very similar numeri- nocal setups, feedbackthe disk evolution and accretion rate of our ‘no- feedback’ case differs somewhat from that described in Stacy primary et al. (2010). Several distinctions between the simulations, however, explain this. The high-density cooling and chem- istry is updated from that used in Stacy et al. 2010 (see Section 2.2). We also use an adaptive softening length in- stead of a single softening length for all gas particles as in secondary Stacy et al. (2010), and our criteria for sink accretion are slightly more stringent. The main contribution to the dif- ference, however, is likely the stochastic nature of the sink particle dynamics. While no sink was ejected in Stacy et al. (2010), the sink ejection and subsequent rapid velocity of the main sink in our ‘no-feedback’ case altered the disk struc- ture, and the main sink would likely have grown to a higher mass otherwise. Nevertheless, the final sink masses in both Figure 7. Projected ionization structure of gas at 1500, 2000, and 3000 yr after initial sink formation. White lines depict the density simulations were still the same to within a factor of two. 40,000 AU Figure 10. Effect of radiative feedback on protostellar accretion. 7.5 8 8.5 9 −3 - a dominant binary has formed (~ 20 and ~10 MO) afterThe radiativ~5,000e feedbac yrk seen here is much stronger than contours of the disk at densities of 10 , 10 , 10 , and 10 cm . Box length is 40,000 AU. Note the pronounced hour-glass morphology Black solid line shows mass growth with no radiative feedback, ii the analytical prediction of McKee & Tan (2008), even when of the developing ultra-compact H region, roughly perpendicular to the disk. This structure gradually expands and dissipates the disk of accretion while black dotted line shows the ‘with-feedback’ case. The blue gas from above and below, reducing the scale height of the disk. solid line is a double powerlaw fit to the sink growth rate for the considering the combined accretion rate of both sinks. They ‘no-feedback’ case, and the red dotted line is a double powerlaw found that a Pop III star could grow to over 100 M⊙ through 4 - Similar mass fitrangefor ‘with-feedbac foundk’ case. in TheHosokawadashed line shows etthe al.growth 2011disk (Science)accretion, as disk shadowing allowed mass to flow onto - after ~5,000 yr: UCHII Region with r~10 AU of the second-most massive sink. Radiative feedback along with the star even while the polar regions became ionized. Our fragmentation-induced starvation leads to a divergence in the ac- lack of resolution prevents this disk shadowing from being cretion histories in less than 1000 yr, and in the ‘with-feedback’ fully modeled on sub-sink scales, and the ionizing photon case the main sink does not grow beyond ∼ 20 M⊙ in the time emission emanating from the sink edge is likely overesti- of the simulation. mated along some lines-of-site. This especially applies to angles just above and below the disk that quickly became ionized in our ‘with-feedback case’, when in reality these cent work by Peters et al. (2010b). They similarly examine regions are unlikely to become ionized until sometime later, ionizing and non-ionizing radiative feedback on massive star only after the gas that is polar the disk becomes ionized first. formation, though they study the case of present-day star McKee & Tan (2008) furthermore assumed disk axisymme- formation, and their initial configuration was different from try, which does not describe the disk in either of our test ours in that they began with a 1000 M⊙ rotating molecular cases. Nevertheless, our shielding prescription does indeed cloud core. They find H ii regions which fluctuate in size and keep the disk from becoming ionized in our ‘with-feedback’ shape as gas flows onto the stars, and that the final mass case. However, the I-front does not expand in a perfectly uni- of the largest stars is set by ‘fragmentation-induced star- form fashion along the polar directions, as the disk rotates, vation,’ a process in which the smaller stars accrete mass and the position of the main sink within the disk varies as flowing through the disk before it is able to reach the most it orbits its companion sink. Thus, different angles will be Figure 8. Projected H2 fraction of gas at 1500, 2000, and 3000 yr after initial sink formation. Gray lines depict the density contours massive star. This is in contrast to models in which the final shielded at different times. Once gas along a certain direction 7.5 8 8.5 9 −3 of the disk at densities of 10 , 10 , 10 , and 10 cm . Box length is 40,000 AU. Note how the morphology of the neutral region is stellar mass is set once ionizing radiation shuts off the disk has been ionized, it may recombine at later times, but LW similar to that of the ionized region. The molecular inflow gradually becomes confined to the disk plane. accretion (e.g. McKee & Tan 2008). In our ‘with-feedback’ radiation prevents most of this gas from cooling back down case we find that fragmentation-induced starvation occurs to below a few thousand Kelvin. The pressure wave of between 1000 and 2000 yr, while afterwards radiative feed- warm neutral gas sourced LW feedback, along with ∼ front is shown in Figure 7, where we can see that the I-front where cs is the ionized gas soundspeed and xs is the position back does lead to a further decline in the sink accretion rate an extended region of neutral gas sourced by for- expands as an hour-glass shape above and below the disk. of the shock in similarity coordinates, which for our case is and the disk mass. Soon after the I-front breaks out from merly ionized particles that recombined, leads to a bubble of neutral gas that continues to expand in all As the I-front expands and widens to encompass a larger xs = 2.54 (see panel b of Figure 9). the sink, the second largest sink accretes much more quickly angular region, it dissipates the more diffuse disk gas above Multiple factors cause deviations from the analytical than the first, intercepting a large portion of infalling mass directions except within a few degrees of the disk plane, and mass flow onto the disk and sinks is greatly reduced. and below the mid-plane, leading to a decline in the disk solution. For instance, the density structure of the gas is that otherwise would be accreted by the main sink. Indeed, the final estimated mass of the main sink is similar to that Thus, while our results underestimate the effect of2 shielding, scale height. Figure 9 compares the I-front evolution with not spherically symmetric. Furthermore, the initially close found in Peters et al. (2010a), where they found the combi- it still highlights how non-axisymmetry will enhance the ef- that predicted by the analytical Shu champagne flow solu- proximity of the ionization front to the sink causes the grav- nation of disk fragmentation and radiative feedback led to a fects of radiative feedback, and the true physical case likely tion (Shu et al. 2002, see also Alvarez et al. 2006). Different ity of the sink to have a non-negligible effect on the early most-massive star of 25 M⊙. However, radiative feedback be- lies somewhere in between our ‘with-feedback’ case and the analytical solutions can be found for the evolution of gas H ii region dynamics, a factor that is neglected in the an- comes important after 2000 yr as it eventually prevents mass prediction of McKee & Tan (2008). In a similar vein, non- under the propagation of an I-front into a powerlaw density alytical solution. However, this effect loses importance as flow onto either sink. A similar study of current-day star axisymmetry can also promote further disk fragmentation, profile, which in our case is approximately ρ r−2. The the H ii region grows well beyond the gravitational radius, formation by Krumholz et al. (2009) found that a prestel- and this in turn can result in N-body dynamics that may ∝ ratio of the un-ionized gas temperature to that of the H ii rg GM∗/cs 50 AU for a 15M⊙ star and 20,000K H ii lar core would similarly collapse into a disk that would host provide another means of reducing Pop III accretion rates. ≃ ≃ region must also be specified for the analytical solution, and region. The analytical solution also does not take into ac- a multiple system of massive stars, even under the effects The fragmentation of primordial gas and growth of in our case we choose a ratio of 20,000 K to 1000 K, or 0.05. count the continued infall of gas onto the central region. of radiation pressure. However, they followed smaller aver- Pop III stars has recently been modeled from cosmological The propagation of the I-front can then be described, assum- The outer envelope beyond the disk, which we take as the ⃝c 2011 RAS, MNRAS 000, 1–18 ing D-type conditions and that the I-front closely follows the gas with density greater than 108 cm−3, grows for the first −2 −1 47 preceeding shock, by a velocity of 2000 yr at a rate of 10 M⊙ yr , or 5 10 neutral particles per second. In≃ comparison, the sink ×typically emits vs = xscs, (15) > 1048 photons per second. Though the infall is not large ∼ ii while the size of the I-front is enough to quench the H region entirely, it is a large enough fraction of the ionization rate to cause the total mass of the rs = xscst, (16) H ii region to fluctuate. The motion of the sink through
⃝c 2011 RAS, MNRAS 000, 1–18 6/26/17
The First Stars: Mass Distribution The Death of the First Stars • Numerical simulations (Heger et al. 2003; Maeder & Meynet 2012 à include rotation) - Bromm, Coppi, & Larson (1999, 2002) - Abel, Bryan, & Norman (2000, 2002) - Nakamura & Umemura (2001, 2002) - Yoshida et al. (2006, 2008); O’Shea & Norman (2007); Gao et al. (2007) - Clark et al. (2011); Greif et al. (2011, 2012); Stacy et al. (2010,12,14,16) Pop I - Hosokawa et al. (2011,16); Latif et al. (2013); Susa (2013); Hirano et al. (2014) • Likely Outcome: Top-heavy initial mass function (IMF)
~ 1 Mo ~ 10-100 Mo dN/dlogM Z Mass
IMF ? PISN Pop I / II Pop III Pop III log M CCSN Initial Stellar Mass
GWs from the Remnants of the First Stars 3
2007; Finn & Chernoff 1993; Cutler & Flanagan 1994), ] 0 our model, Pop III
-3 10 our model, Pop I/II ∞ 2 -1 2 |h(f)| 10 Kinugawa+16, Pop III
ρ =4 df, (1) Mpc -2 Belczynski+16, Pop I/II !0 Sn(f) -1 10
yr -3
where h(f) is the Fourier-domain (sky- and orientation- ⊙ 10 -4 averaged) GW strain at the detector, and Sn is the 10 noise power spectral density of a single aLIGO de- 10-5 SFR [M tector. We assume that an event is detectable if 10-6 ρ > 8, as conventionally done in the LIGO literature 0 0.2 0.4 0.6 0.8 1 (Abadie et al. 2010; Dominik et al. 2015; Belczynski et al. cosmic time [Gyr] 2016; deMink&Mandel2016). (This translates to SNR larger than 12 for a three-detector network, e.g. the two aLI- Figure 1. Comparison of our SFR for the fiducial IMF with the GOs and advanced Virgo.) For the current aLIGO detectors, models used by Kinugawa et al. (2016)(hereafterK16)andB16. OurPopI/IISFR,whichisadoptedfromBehroozi & Silk (2015), Gravitationalwe Wave use the O1(GW) noise power Archaeology spectral density (Abbott et al. GWs from the Remnants of the First Stars is in good agreement with the corresponding SFR by B16.Forthe 2015),whereastoassessdetectabilitywhenthedetectorsare PopIIIstars,ourself-consistentmodellingyieldsapeakSFRthat in their final design configurations we use the zero-detuning, is about an order of magnitude lower than the value by K16. high-power configuration of Abbott et al. (2009). For h(f), Downloaded from weuseeitherinspiral-only,restrictedpost-Newtonianwave- Merger Rate vs. Cosmic time • LIGO GW150914: forms (computing the Fourier transform with the sta- tionary phase approximation, see e.g. Maggiore 2007), or 2
• LIGO: GW150914 ])
inspiral-merger-ringdown PhenomA (non-spinning) wave- -3 forms (Ajith et al. 2008, 2009). We employ the former for 1 - Pop III BH-BH could http://mnrasl.oxfordjournals.org/ Gpc
BH-NS, and NS-NS systems, with a cut-off at the frequency -1 possibly provide source - inspiral and merger of 0 of the innermost stable circular orbit (ISCO; note that (see also: Belczynski et al. 2016) the ISCO frequency BH-BH also corresponds binary with approximately to the -1 merger frequency of NS-NS systems). For BH-BH systems, à M1=36 M¤ -2 particularly at high masses, the merger-ringdown contains - But: Multiple other à M =29 M¤ considerable SNR, hence2 we use PhenomA waveforms. We -3 IMF 1->100 scenarios proposed! then calculate the detection rate as (Haehnelt 1994) IMF 3->300 -4 IMF 10->1000 2 2 Kinugawa+16-IMF 10->100 dn - Big Q: Whered R diddt systemdL =4πc dzdm1dm2 , (2) Kinugawa+16, rescaled - Break degeneracy with
log(Merger Rate Density [yr -5 dt ρ>8 dm1dm2 dz 1+z Dominik+13, 0.1Z⊙ ! " # at University of Texas Austin on May 10, 2016 z
frequency at the detector). Our model’s prediction for −3 −1 −3 with a peak value of SFRmax =3× 10 M⊙ yr Mpc , ΩGW should be compared with the 1σ power-law inte- grated curves (Thrane & Romano 2013)inAbbott et al. which is about an order of magnitude higher than our re- (2016a)fortheaLIGO/advancedVirgonetworkintheob- sult. The SFR is about the same for all our Pop III IMFs, serving runs O1 (2015-16) and O5 (2020-2022), which repre- because we calibrate each model to match τ.Theexactred- sent the network’s sensitivity to standard cross-correlation shift evolution of the Pop III SFR depends on the details searches (Allen & Romano 1999)ofpower-lawbackgrounds. of the treatment of reionisation and metal enrichment (cf. Johnson et al. 2013, for a comparison). Theintrinsicmergerratedensityofcompactobjectscan 3RESULTS be seen in Fig. 2. To compare with K16 we rescale their SFR to our peak value, and run a case with their IMF (10-100). In Fig. 1, we compare our star formation rate (SFR) to Our values are about an order of magnitude lower than the 3 other models. Our self-consistent Pop III SFR, with a peak rescaled prediction by K16 in the regime relevant for GW −4 −1 −3 value of SFRmax =2× 10 M⊙ yr Mpc , is in com- detections, i.e. mergers occurring at late cosmic times. This pliance with Visbal et al. (2015), who show that it can- is to be ascribed to the different binary evolution models. notexceedafew×10−4,withoutviolatingtheconstraints Given that K14 study explicitly Pop III star binary evolu- set by Planck Collaboration (2015). K14 assume an SFR tion, and we extrapolate models at higher metallicity, we 6/26/17
BH Feedback and Early Galaxy Formation First Galaxies: Assembly with BH Feedback (Jeon, Pawlik, Bromm, & Milosavljevic 2014, MNRAS, 440, 3778)
10 Jeon et al. no HMXB with HMXB - Input of heat, pervasive (partial) ionization - Feedback: positive or negative?
Gas density - Vigourosly explored, e.g.: - Ricotti & Ostriker 2004; Kuhlen & Madau 2005; Zaroubi et al. 2007; Alvarez et al. 2009; Mirabel et al. 2011; Park & Ricotti 2011
Di Matteo et al. 2005 Gas temperature
Figure 5. Hydrogen number densities (top)andtemperatures(! bottom)atz =19.6, averaged along the line of sight, in a cubical cutout • High-Massof linear X-ray extent 300 comovingBinary kpc centred (HMXB) on the first sink particle,Mirabel in the simulationset al. (2011); without (left)andwith( Fragosright et)HMXBs.The al. (2013) stronger photo-ionization heating in the presence of HMXBs and the associated increase in the thermal gas pressure imply a stronger smoothing of the gas density field.
that included persistent X-ray radiation from an isolated BH n¯H is the cosmological mean gas density, and fesc is the miniquasar accreting di↵use halo gas and that started from fraction of ionizing photons that escape haloes to ionize the identical initial conditions. We identify two key di↵erences IGM (e.g., Razoumov & Sommer-Larsen 2006; Gnedin et al. in the simulation methodology to explain this di↵erence in 2008; Wise & Cen 2009; Paardekooper et al. 2011; Yajima star formation rates. First, in the current work, we account et al. 2011). Figure 4 shows that the simulated star forma- for shielding of the primordial molecular gas from LW pho- tion rate densities are generally below the critical star forma- tons that was ignored in the previous simulations. LW feed- tion rate density, indicating that the reionization of the sim- back is therefore less e↵ective in the present work than in the ulated high-resolution region remains incomplete. Because, previous simulations and this promotes gas cooling and con- typically, C>1(seeSection3.3)and,bydefinition,fesc 6 1, densation. Second, in Jeon et al. (2012), we overestimated we have conservatively underestimated the ratios of critical the radii of HII regions around Pop III stars by a factor of and simulated star formation rate densities. 2 as a result of a numerical error discussed in Jeon et al. ⇠ (2012). Thus, in the current work, the suppression of star formation by stellar ionizing radiative feedback is expected 3.3 E↵ects of HMXBs on IGM properties to be weaker than in the preceding paper. Figure 5 shows images of the gas density and temperature It is interesting to compare the simulated star formation in the high-resolution region centred on the site of forma- rate densities with the comoving critical star formation rate tion of the first stellar binary near the end of the simulation density to sustain ionization against recombinations (Madau at z =19.6. In the simulation that included ionization by et al. 1999), HMXBs, the gas is photoheated to temperatures in excess 3 of a few times 103 K throughout the simulation volume. 1 3 C 1+z ⇠ ⇢SFR,crit =0.02 M yr Mpc , (24) In the simulation without HMXBs, which misses the bright fesc 20 ✓ ◆✓ ◆ HMXB phases of the accreting BHs, on the other hand, the where C n2 /n¯2 is the clumping factor that average temperature is significantly lower, and heating to ⌘hHi H parametrizes the recombination rate in the IGM (e.g., comparably high temperatures is limited to volumes inside Madau et al. 1999; Miralda-Escud´eet al. 2000; Wise & Abel stellar HII regions. The higher temperatures in the simu- 2005; Pawlik et al. 2009; Finlator et al. 2012; Shull et al. lation with HMXBs imply a cosmological Jeans mass (e.g., 2012), the brackets indicate the volume-weighted average, Shapiro et al. 1994; Gnedin 2000; Wise & Abel 2008a) and
c xxxx RAS, MNRAS 000,1–??
Assembly of the First Galaxies Assembly of the First Galaxies (Jeon, Bromm, Pawlik & Milosavljevic 2015, MNRAS, 452, 1152) (Jeon, Bromm, Pawlik & Milosavljevic 2015, MNRAS, 452, 1152) - SPH/TRAPHIC, cosmological ICs Pop III/NS - SPH/TRAPHIC, - galaxy collapses
cosmological ICs Pop II at z~10 Pop III/BH - galaxy collapses
at z~10 - Self-consistent feedback à One-star-at-a-time - Self-consistent feedback
à One-star-at-a-time - Ubiquitous metallicity floor à First galaxies are pre- dominantly Pop II systems
4 6/26/17
THE FIRST GALAXIES 9
virial temperature of the hosting halo with a virial mass Mvir: M 2/3 1+z T 104 K vir . (18) vir ⇥ 5 107 10 ⇥ ⇥ 4 Above Tvir 10 K, the gas within the halos is able to cool mainly via atomic hydrogen. In Figure 6, we show the evolution of the virial mass for the most massive halo, which will host the first galaxy at z 10, as well as of the total and cold gas masses for the⇥ three simulations presented in this work. The virial mass of the DM halo is estimated as the mass within a radius where the average BH Feedback: Suppression of SeedDM density Growth is 200 0(z ), where 0 is the mean cosmic Pathways to Early SMBH Formation density at a given⇥ redshift. Cold gas is identified by the (Jeon, Pawlik, Greif, Glover, Bromm, Milosavljevic &condition Klessen that 2012 the gas, temperature is less than half of THE FIRST GALAXIES 9 the halo virial temperature, T<0.5 Tvir(z). We also 2 A. SmithBH et al. Mass vs. Cosmic Time indicate the critical ApJ mass, 754, required 34) for the onset of atomic virialhydrogen temperature cooling at of a the given hosting redshift. halo with a virial mass M : Wevir find that at z 18, the halo is dominated by hot yield a near-extremal Kerr black hole if cooling does not initiate gas, exceeding the amount of cold2/3 gas by an order of 4 Mvir 1+z magnitude.Tvir As10 timeK passes on, the cold gas. mass(18) in- fragmentation during the collapse (Rees 1984). Accretion creases, eventually⇥ accounting5 107 for 80%10 of the total gas - - - Eddington -limited ⇥ ⇥ Early simulations of collapsing primordial gas clouds showed mass in simulations4 BHN and BHS. This trend can be Rate understoodAbove Tvir by the10 vulnerabilityK, the gas within of the the halo halos gas isto able stellar to that most of the gas does indeed fragment into dense stellar clumps coolradiative rate mainly feedback. via atomic The hydrogen. corresponding In Figure evacuation 6, we of show gas which eventually virialize into a spheroidal galactic bulge (Loeb & thefrom evolution the halo of is the very virial strong mass at high for the redshifts, most massivez 18, halo, be- which will host the first galaxy at z 10, as well as of Rasio 1994). The emergence of a SMBH would thus be forestalled. cause the halo potential wells were not⇥ yet deep enough theto retain total photo-heated and cold gas gas. masses We for note the that three the simulations sharp dips On the other hand, if a central seed black hole of mass & 106 M presentedin the amount in this of work. cold gas The are virial due mass to star of the formation DM halo in- is estimatedside the halo as the itself mass (at withinz 12), a radiusas well where as in neighboring the average were to form during the collapse then it could quickly grow by DM density is 200 (⇥z), where is the mean cosmic BH Mass halos that are su⇥⇥ciently0 close, within0 ⇥ 5 kpc. The ion- steady accretion to a quasar size black hole (Li et al. 2007). Oth- densityizing feedback at a given from redshift. the accreting Cold gasBHis starts identified operating by the at conditionz<18, above that which the gas the temperature e ect is too is weak less as than a result half of erwise dynamical instabilities in the inner region of the disk would the initially halo virial very temperature, low BH accretionT<0 rates..5 Tvir(z). We also inhibit the accretion required for such growth. Thus, the crucial indicate the critical mass required for the onset of atomic While the total amount of gas is not sensitive to the bottleneck is the formation of the initial, million solar-mass, seed. hydrogenBH feedback, cooling the at reduction a given redshift. in cold gas mass by a fac- torWe of find 5 indicates that at thatz 18, the the additional halo is dominated heating from by this hot Either way, realistic modeling of massive black hole inception is gas,feedback, exceeding on top the of the amount stellar of feedback, cold gas has by a an significant order of highly complex and ultimately relies on understanding both small- magnitude.impact on the As gas time in the passes center on, of the forming cold gas galaxy. mass in-As - NegligibleFig. growth 8.— BH growth of central with and without BH feedback.when Shownit is isradiatively the creases, active eventually accounting for 80% of the total gas and large-scale phenomena. For example, it was pointed out that redshift evolution of the BH accretion rate, the temperature and the halo grows further via smooth accretion and mergers masswith minihalos, in simulations however, BHN at andz BHS.13, bothThis trendthe total can gas be (Smith, Bromm & Loeb 2017, A&G, in press; arXiv:1703.03083) density of the gas in the immediate vicinity of the BH, as well as the understood by the vulnerability⇥ of the halo gas to stellar low angular momentum configurations would help facilitate run- - Significantlyresulting BHsub- massesEddington for simulations BHN (black) and BHS (red ). mass and mass of cold gas are no longer sensitive to the Figure 1. Schematic diagram illustrating the timing problem for supermas- In the top panel, we also indicate the corresponding Eddington- radiativeBH radiative feedback. feedback. The At correspondingz 11.5, the evacuation condition forof gas an away collapse due to the significantly lower centrifugal barrier; in limited accretion rates for the two cases (dashed lines). fromatomic the cooling halo is halo very is strong satisfied. at⇥ high redshifts, z 18, be- sive black hole growth in the early Universe. The observed quasars with - A “head-start” scenario:9 10 Direct collapse black holes (DCBHs) this context, quasar seeds could be a natural consequence of the ini- - Need for more exotic pathway à direct-collapse BH? cause For simulation the halo potential BHB, the wells heating were from not yet the deep HMXB enough is so a mass M 10 M at redshift z 6 7 could have originated from BHN and BHS cases, by a factor of about 3. This is a to retain photo-heated gas. We note that the sharp dips • ⇡ ⇡ tial collapse of regions with unusually small rotation (Eisenstein & (Bromm & Loeb 2003; review of recent developments àstrong Volonteri that the & meanBellovary temperature 2012) of the gas inside the massive seeds or alternatively from low-mass stellar progenitors with hyper- consequence of the positive feedback, where gas collapse inhalo theT> amount0.5 T of( coldz), over gas theare dueentire to range star formation of simulated in- Loeb 1995). Such low-spin cosmological perturbations might pro- into distant minihalos is facilitated via H cooling pro- side the halo itselfvir (at z 12), as well as in neighboring Eddington growth rates. 2 redshifts z 15, and the⇥ total gas mass is reduced by an vide environments for which SMBH formation is an extreme man- moted by the strong X-ray emission from the HMXB. As halosorder that of magnitude are su⇥ciently by photo-evaporation, close, within ⇥ 5 kpc. as is The clearly ion- density fluctuations grow, however, and hence star for- izingevident feedback in the from dot-dashed the accreting linesFigure BH starts 6. operating This result at ifestation of the ⇤CDM model. mation is enhanced, the local negative radiative feedback z
tion pathways include stellar-remnant black holes that have under- & Natarajan 2014). outflow gone vigorous gas accretion, dense star clusters in which runawayIGM On the other hand, forming a DCBH requires collapse with- Metal collisions trigger the formation of a SMBH, or a cluster of post- out fragmentation, a cosmic “miracle” of sorts that is naturally ex- supernovae neutron stars or black holes coalescing in the dynam- plained within the context of galaxy formation theory. The idea is
time enrichment ically unstable central core. In the context of(Becerra cosmology, et al. SMBHs 2017; arXiv:1702.03941)that if a primordial gas cloud, devoid of any heavy chemical ele- can grow over longer timescales via episodicFigure 3. A galaxycartoon illustration mergers of the galactic and outflow modelments for CR7. (“metals” A in astronomical terminology), is bathed in a sea strong source of radiation drives an expanding shell of gas from the centre accretion from streams of cold gas alongthrough the filaments interstellar medium of (ISM). the Thecos- Ly↵ photons areof redshifted ultraviolet by (UV) radiation then it will be unable to cool and the outflow and then transmitted through the intergalactic medium (IGM) (Smith etmic al. web2017; (Mayer arXiv:1703.03083) et al. 2010, 2015). Still,which it lies was between recognized the distant galaxy that and the a observer. Theform combined stars ra- (see Fig. 2). Specifically, according to the Jeans crite- Figure 2. Schematic density-temperature relation for the collapse of pri- diative transfer through the ISM and the IGM determines the observed Ly↵ 6 mordial gas ingas an atomic cloud cooling may halo conceivably with a virial temperature bypassT conventional star formation and rion for triggering gravitational instability, high thermal pressure ~10 M first stars vir ⇠ line profile. z~20 o 104 K. The evolutionary track in the pressence of a strong LW flux can be significantly di↵erent from the typical scenario with H2 and metal line cooling. MNRAS 000, 1–6 (2017) radiates broadband emission from optical to X-ray wavelengths - DCBH formation: Strong LW flux suppresseswith a peak in the SF UV. Theduring characteristic collapse blackbody temper-of metal- Dark matter halos is required to prevent the gas cloud from fragmenting. In present- ature for the Eddington luminosity near the event horizon is 5 8 1/4 TEdd 85 10 K(M /10 M ) (Rees 1984). DCBHs are freeday star gas forming into clouds, atomic line cooling by cooling heavy elements halo and dust (Mvir~10⇡ M⇥ ¤) • e ciently radiates away thermal energy that would otherwise pro- born in gas rich environments and may be self-shielding to ion- vide stability. However, in the early Universe only hydrogen and izing photons. Still, the gas remains transparent to X-rays and the helium were available to cool the gas. Thus, a small abundance of “Compton-thick” spectrum retains the non-thermal tail contributed by Bremsstrahlung radiation and magneto-hydrodynamical (MHD) molecular hydrogen (H2) played a key role as the primary cool- ing agent below the 8, 000 K accessible via atomic hydrogen processes (Pacucci et al. 2015b). In these environments, radiation ⇠ line cooling. If strong non-ionizing Lyman-Werner radiation (LW; pressure from the black hole along with concurrent star forma- with photon energies below 13.6 eV) from neighboring galaxies tion and nearby supernovae likely had a substantial impact on the photodissociates H then the evolutionary track through density- host galaxy (Jeon et al. 2012). Lower mass minihaloes with shal- 2 5 7 lower gravitational potential wells (Mvir . 10 M ) would have temperature phase space is significantly altered (see Fig. 2). Only atomic cooling primarily through the Lyman ↵ (Ly↵) line of hydro- been especially susceptible to radiative feedback, which potentially gen e ciently keeps the collapsing core below 104 K. The first depleted the reservoir of gas needed to fuel black hole growth ⇠ (Whalen et al. 2004; Wise et al. 2012). The viability of the DCBH idealized models of isothermal collapse including H2 photodisso- ciation still predicted the cloud breaks up at late stages despite the mechanism also relies on the regulation of chemical feedback be- higher temperature track (Omukai 2001). Gas collapse may be en- cause molecular hydrogen and metal cooling induces fragmenta- tirely suppressed in the minihaloes, with total (or virial) masses tion (Hartwig et al. 2016, recall Fig. 2). Therefore, the eventual 6 DCBH formation sites must remain free from star formation dur- of Mvir . 10 M , that are predicted to host the first stars, when the LW background flux is above a critical value. More massive ing collapse. Furthermore, the emission of UV and X-ray photons host haloes, on the other hand, in general become self-shielding so promotes H2 formation and increases the critical LW flux needed that molecules can form again, eventually leading to star formation to form DCBHs (Inayoshi & Tanaka 2015; Latif et al. 2015). (Oh & Haiman 2002). However, Bromm & Loeb (2003), carrying A particularly interesting source of feedback in DCBH envi- out the first simulations of this collapse with cosmological initial ronments is Ly↵ photon trapping. In the vicinity of a newly col- conditions, recognized that under conditions of unusually strong lapsing primordial gas cloud, the surrounding neutral hydrogen is LW irradiation, the inflow would continue on its near-isothermal extremely optically thick to photons near the Ly↵ resonant line. In track. The resulting free-fall collapse of the atomically-cooling gas one-dimensional simulations, a dense shell-like outflow structure could then produce massive black holes directly, without under- forms in hydrodynamical response to the central source ionizing going any sub-fragmentation and star formation. Subsequently, the and heating of the gas. As illustrated in Fig. 3, the Ly↵ photons are DCBH model has received increased attention (e.g. Begelman et al. redshifted by the galactic outflow, producing a velocity o↵set that 2006; Regan & Haehnelt 2009; Choi et al. 2013, 2015). This body may result in considerably less scattering out of the line of sight by of work is broadly reviewed in Volonteri (2012), Haiman (2013), the intervening intergalactic medium (IGM). Due to the complex Loeb & Furlanetto (2013), Johnson & Haardt (2016), and Latif & nature of the Ly↵ radiative transfer, explorations of Ly↵ feedback Ferrara (2016). tend to focus on order of magnitude estimates based on idealized calculations (Oh & Haiman 2002; McKee & Tan 2008; Milosavl- 5 jevic´ et al. 2009). With the aid of post-processing Monte-Carlo ra- diative transfer (MCRT) methods, Dijkstra & Loeb (2008) found Radiation from the first black holes that multiple scattering within high H i column density shells is ca- Active galactic nuclei are conspicuous manifestations of gas ac- pable of enhancing the e↵ective Ly↵ force by one or two orders of cretion onto supermassive black holes, with luminosities exceed- magnitude. The first self-consistent Ly↵ radiation-hydrodynamics ing the total starlight of their host galaxies. The accretion disk (RHD) simulations were performed by Smith et al. (2017), who
MNRAS 000, 1–6 (2017) 6/26/17
Are First Galaxies really all enriched already? The curious case of CR7 ! The curious case of “COSMOS Redshift 7” (CR7) ! Not Pop III Stars but “Direct Collapse BH”!? (Sobral, Matthee, et al. 2015, ApJ, 808, 139) - Hartwig et al. 2016, MNRAS, 462, 2184 The Astrophysical Journal, 808:139 (14pp), 2015 August 1 Sobral14 et al. T. Hartwig et al.
12 11.2 PopIII stars 1σ PopIII stars 3σ - Need to reconcile high - Luminous Lyα emitter at z=6.6 PopIII BH 1σ PopIII BH 3σ 11 DCBH 1σ DCBH 3σ 11.5 44 -1 -1 luminosity with “metal tax” (LLyα≈10 erg s ) 10.8 -3 10.6 (=assembly of massive source ) 11 ) -2
⊙ - Also: strong HeII 1640Å ⊙
10.4 ) -4 /M /M gas will inevitably imply enrichment) ) A 10.5 43 -1 10.2 D
-3 ⊙ (L1640≈10 erg s ) /M 10 -5 log(M log(M
metal - DCBH plausible candidate 10 5 => 9.8 -4 Teff ~ 10 K log(Z/Z -6
9.6 log(M for ~metal-free, massive system 9.5 - 3-component system: -5 9.4 -7 => Observational agenda: ---9 B+C: enriched (Pop II) 9.2 CR7 9 9.5 10 10.5 11 11.5 12 -6 0 1 2 3 4 5 6 7 - deeper NIR spectra --- A: only log(MupperC/M limits⊙) on Z log(mass/M⊙) => Pop III ??? - radio, X-ray Figure 15. Masses of the second, third, and forth most massive Figure 16. Comparison of the three different models, regarding Figure 5. Left: the SED of “CR7,” from observed optical (rest-frame FUV) to observed MIR (rest-frame optical) and the best fit with a normal stellar populationhaloes(not at z =6.6, assuming that they merge to one halo within the metallicity of clump A at z =6.6andthemassofthestellar. including PopIII stars). Red crosses indicate the flux predicted for each broadband filter for the best fit. The fit fails to reproduce the strong Lyα emission line and100 also Myr. The grey region indicates the required mass range for population(see or of the also: BH, respectively. Agarwal The et contours al 2016; illustrate Pallottini the et al. 2015; Smidt et al. 2016) the excess in J band, due to He II emission. Moreover, and even though the fit is unable to reproduce all the information available for the source, it requires anclump age of A and C of CR7 and we show 512 different merger tree probability of 68% (1σ)and99%(3σ)tofindanobjectwithin −1 10.3 700 Myr (the age of the Universe is 800 Myr at z = 6.6). In this case, the galaxy would have an SFR of ∼25 Me yr and a stellar mass of ∼10 Me. This, of course,realisations. The fourth most massive clump D should be signifi- this region in the mass-metallicity-plane. The grey shaded area is not able to explain the strong Lyα and the strong He II emission line. Right: same observed SED of CR7 as in the left panel plus HST photometry for clumps B+C (magenta triangles). The black line shows a fit with a pure PopIII SED to the rest-frame UV part; the magenta line the SED of an old simple stellar population withcantly 1/5 less massive than clump C, because otherwise one should indicates the region in which we expect the system to have the solar metallicity which matches the flux from clumps B+C; the green line shows the predicted SED summing the two populations after rescaling the PopIII SEDsee by a a fourth clump present in CR7. In this sample, we find con- same observational signature as CR7. A Pop III stellar population 7 factor of 0.8. The composite SED reproduces well the observed photometry. Although there is tension between the strength of the He II line and nebular continuumstellations in the right mass range and with MD/MC ! 0.1. The has to have a mass of ∼ 10 M⊙ and a black hole should be in the fi 12 5 6 emission (cf. text), a PopIII contribution is required to explain the He II λ1640 line and the corresponding excess in the J band. Although He II is very strong, wemostnd massive clump B has always a mass of ∼ 10 M⊙. mass range 10 − 10 M⊙ to yield the observed line luminosities. no evidence for any other emission lines that would be characteristic of an AGN. Furthermore, the clear IRAC detections and colors, and particularly when taken as a The upper limit of the metallicity is set by Z ,inordernot whole, can be fully explained by a PopIII population, while a normal stellar population or an AGN is simply not able to. limit to violate the observed metal line ratios. BHs of different origin In contrast, the DCBHs reside by construction in low- might reside in metal-poor halos at z =6.6, but only DCBH are formation rate (SFR) or bursts are explored in this case. We much for different star formation histories (SFHs): bothmetallicity for haloes with a much higher gas mass of up to massive enough to account for the observations. also explore SEDs from “composite” stellar populations, exponentially declining and rising cases, the fits prefer long 9 ∼ 10 M⊙.Thesehighermassescanbereachedbecause showing a superposition/mix of PopIII and more normal timescales (i.e., slowly varying SFHs). Whereas for exponen-we quench Pop III star formation in the host halo with a populations. tially declining SFHs and for constant SFR the best-fit halo D should be at least an order of magnitude less mas- strong LW flux from a nearby star-forming halo. We find Nebular emission from continuum processes and emission attenuation is negligible (A = 0), a higher attenuation is sive than halo A (M /M ! 0.1). This mass distribution is V that DCBHs reside in a metal-free environment for only sev- D A lines are added to the spectra predicted as described in Schaerer needed for rising SFHs, as expected (cf. Schaerer & Pelló 2005; the anticipated constellation of CR7. We analyse the metal- eral tens of Myr before they are polluted by merging with the & de Barros (2009). Nebular emission from continuum Finlator et al. 2007). Depending on the assumed metallicity, licity, and estimate the masses of both Pop III stars (with neighbour halo. DCBHs therefore reside in low-metallicity processes and emission lines are proportional to the Lyman this may reach from AV = 0.5 (for 1/5 solar) to 0.25 (for 1/200 the fiducial model) and BHs resulting from the two differ- – −1 haloes only down to z ≈ 8.5, which is ∼ 240 Myr before continuum photon production. Whereas many emission lines Z ). The corresponding SFR is ∼30 40 Me yr , and the ent seeding scenarios in halo A at z =6.6. The results are are included in general, we only include H and He lines for the predicted IR luminosity ranges from log(L /L )=10.4z =6 to .6. When interpreting this result, one should keep in IR shown in Figure 16. For a given scenario, this plot illustrates PopIII case (cf. Schaerer 2003). The IGM is treated with the 11.2, assuming energy conservation. The typical specificmind SFR, that the rapidly-changing metallicity of the host halo −1 the probability for a certain mass-metallicity combination. prescription of Madau (1995). Attenuation by dust is described sSFR, obtained from these fits is sSFR ∼ 1.2–1.9 Gyr due, as to the merger is an artefact of our recipe for DCBH “ ” For Pop III stars, the stellar mass in clump A at z =6.6 by the Calzetti law (Calzetti et al. 2000). expected for a mature stellar population with SFH closeformation. to For this analysis, we use an optimistic value of 3 constant c.f. González et al. 2014 . remains always below 10 M⊙ and the halo has a metallic- ( ) Jcrit = 100. The actual value can be up to an order of mag- ity of ∼ 10−3.PopIIIremnantblackholesmightresidein 5.1. CR7: SED Fitting with a Normal Population nitude higher (Section 5) so the derived values should be metal-poor halos at z =6.6, but are not massive enough 5.2. CR7: SED Fitting with Contribution from PopIII Starstreated with caution, since a higher Jcrit might inhibit the 3 Part of the photometry of CR7 is explained relatively well (! 10 M⊙)toexplainCR7.OnlytheDCBHscenariocan with “standard” models, as illustrated in Figure 5 (left panel). The presence of strong Lyα and He II emission lines, plusformation the of DCBHs. A lower value of Jcrit may lead also to 4 5 explain the observed line luminosities. For the given mass fi lighter seeds with masses in the range 10 − 10 M⊙. The shortcomings of these ts is that they cannot account for absence of other UV metal emission lines (cf. above), may be constellation of the halos A, B, C, and D, the probability the relative excess in the J band with respect to Y, H, and K, due to exceptionally hot stars with a strong and hard ionizing that clump A hosts Pop III stars is 37%, that it hosts a and that the Lyα emission is not strong enough to reproduce flux, resembling that expected for PopIII stars (cf. Tumlinson Pop III remnant BH is 98% and that it hosts a DCBH is the entire flux observed in the NB921 filter. Both of these et al. 2001; Schaerer 2002).Afit with PopIII templates (4.6from Comparison of scenarios shortcomings thus relate to the presence of the strong emission Schaerer 2002, 2003) is shown in Figure 5 (right panel: black 10%. These probabilities add up to over 100%, because due lines (Lyα and He II) observed in spectroscopy and also line). PopIII models (SchaererThe2002, 2003curious) show that at 3.6caseInμ thism section, of weCR7 compare the three models and estimate, to mergingPerspectives: of the progenitor halos, clump A can contain affecting the broadband photometry. (for z = 6.6) there are strong He II and He I lines, apartwhich from has the highest probability to reproduce the observa- Pop III stars and black holes at the same time. Compar- The typical physical parameters derived from these SED fits Hβ, and that these! Not Helium Pop emission III lines Stars should havebutfltions uxes“Direct of CR7. Collapse To do so, we assume BH”!? that CR7 is an ongoing ing these results to, e.g., Figure 11 shows that it is crucial which only include “normal” stellar populations indicate a comparable to that of Hβ in the case of PopIII (Schaerer 2002merger). of the three clumps A, B, and C, which merge at to treat CR7 as three individual clumps. The third most 10 −1 fi stellar mass M 2 10 Me, SFR ∼ 25 Me yr , and a Speci- Smith,cally, 3.6Brommμm should & be Loeb contaminated 2016 by, HeMNRAS,II 4686,z 0 He460,=6I .31430. This is the most plausible scenario to obtain the massive halo A is less affected by metal enrichment, but the * 3 fairly old age (∼700 Myr; the universe is ∼800 Myr old at 4471 and He I 5016. For 4.5 μm, apart from Hα, the He I right5876 constellation of masses in CR7 as shown in Figure 15. Pop III remnant BHs do not merge to masses above 10 M⊙. z = 6.6). The SED fits and the derived parameters do not vary should be detected and could be comparable to Hβ. The general The most6 massiveA. Smith halo et al. B always fulfils the requirement of Under these conditions and with the optimistic assumption 12 Flux vs. velocity offset MB ≈ 10 M⊙,whereasonly< 10% of the merger tree of Jcrit =100wefindthat∼ 90% of the DCBHs reside in 8 realisations yield masses in the right range for halo A and suffi•ciently Very metal-poor dynamic, gas at z rapidly=6.6. developing C. Out of this subset, a small fraction also meets the crite- Only ∼ 0.5% of all the merger tree realisations, yield the Lyα rion that the fourth most massive halo D is significantly less rightfield mass constellation and out of these, only 9% produce massive than the third most massive halo. We assume that results that are consistent with the observations of CR7. Ac-
outflow • Closing the final⃝c 2015 gap RAS, in MNRAS our worldview000,1–19
• Driven by supercomputers and our best telescopes
• The high-redshift frontier: How did it all begin?
Reviews: Figure 4. Observed line-of-sight flux as a function of Doppler velocity Figure 6. Evolution of the radial surface brightness profile for the Pop III v = c / for the Pop III and MBH models. The grey curves are intrinsic and MBH models at t = 1, 2, 3, 5, 8.5, 60 Myr. The late dashed curve • { Bromm} & Yoshida 2011, ARA&A, 49, 373 - With DCBH central source, can sufficientlyprofiles while accelerate the black and red curves represent shell a plausible reprocessing is from a simulation with lower spatial resolution. For the intrinsic grey due to scattering in the IGM based on the model from Laursen et al. (2011) curves the region is more uniform and extended in the Pop III case while red • Karlsson, Bromm, & Bland-Hawthorn 2013, Rev. Mod. Phys., 85, 809 and ⌧GP from Equ. (6), respectively. These are shown for early to late times the MBH emission more closely follows the shell expansion as a result of of 0.5 Myr (dotted), 3 Myr (solid), and 8 Myr (dashed). The observed veloc- increased Ly↵ trapping. The black and red curves are corrected for trans- 1 • Bromm 2013, Rep. Prog. Phys., 76, 112901 ity o↵set of +160 km s (vertical line) is roughly reproduced in the MBH mission through the IGM based on the model from Laursen et al. (2011) red (see also: Dijkstra, Gronke & Sobral 2016)case but not in the Pop III model. The uncertainty due to DEIMOS/Keck and ⌧GP from Equ. (6), respectively. After scattering in the IGM the Pop III and X-SHOOTER/VLT spectral resolution is illustrated by the grey vertical model appears much more• compact Smith, whereas the MBH Bromm, model remains spa- & Loeb 2017, A&G, in press (arXiv:1703.03083) regions (see Section 2.1 for additional details). tially extended.
low vsh to be a few times larger, we nonetheless find it less probable that the CR7 source is significantly younger than this. Again, such long source durations can more easily be explained within a MBH scenario than a Pop III starburst.
3.4 Caveats We now discuss the limitations of our simulations and their poten- tial impact on the above results. Uncertainties in the galaxy model parameters and initial conditions may produce slightly di↵erent ob- servables. For example, normalizing the black hole flux to the ob- served He ii emission leads to a much lower Ly↵ luminosity. For this reason we have explored the “Boost” and “Leak” models to Figure 5. The time evolution of the velocity o↵set for the red peak of the determine the impact of extra Ly↵ emission or less ionizing ra- intrinsic line-of-sight flux. The Pop III and MBH models are shown by diation. We also modeled the photon emission as a central point the solid black and red curves, while the Boost modification is denoted by source. While the Pop III starburst is initially centrally dominated, dashed curves of the corresponding color. The location of the o↵set is fairly constant throughout the simulations. it is highly unlikely to remain so as star formation proceeds. How- ever, similar to Dijkstra & Loeb (2008) we expect our final results are not particularly sensitive to the choice for the radial Ly↵ photon a slight decrease in the velocity o↵set at larger radii which is ex- emission. plained by the resonant scattering bias to return to line-centre. The Another issue is that of numerical resolution. At a halo mass primary di↵erence between the early and late times is the radius of M 3 1010 M we achieve Lagrangian mass elements of vir ⇡ ⇥ of the shell front which by 60 Myr has advanced to 4 kpc and 6.6 105 M which at various times corresponds to 0.1 2 pc ⇠ ⇡ ⇥ ⇡ 8 kpc in the MBH and MBH Boost models, respectively. At later resolution at the shell front. Improving the resolution by a factor ⇠ times the velocity o↵set decreases, possibly as a result of the lower of 5 produces very similar results except at the earliest times or ⇡ shell velocity and a more di↵use Ly↵ trapping environment. Still, innermost radii. Higher resolution leads to slightly faster initial ve- all of the aforementioned Ly↵ features persist in simulations carried locities because higher densities allow for more e cient radiation 6 out to 120 Myr. If the shell radius is indicative of the Ly↵ emitting coupling. However, the computational cost may change dramati- 1 region then even at v 100 km s it would take at least 50 Myr cally based on the number of timesteps with MCRT calculations sh ⇠ to achieve a size of 10 kpc in diameter. Although our models al- and the convergence criteria. To be safe we require that between ⇠
MNRAS 000, 1–9 (2016) 6/26/17
Standard Model for Pop III SF: Sensitivity to Small-scale Nature of DM
• (Ubiquitous) Disk Fragmentation à small multiple FDM à Absence of high-z Post-starbursts 2nd gen. stars . z~10 8 M ~ M / t ~ T 3/2 (Pop I: T ~ 10 K, ~10 Mo first galaxy J ff st . Pop III: T ~ 300 K) 1 low-M stars? M
feedback time
6 ~10 Mo first stars (dust-free à reduced opacity) Core: z~20
M~MJ Dark matter halos
7