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10.1146/annurev.astro.42.053102.134034

Annu. Rev. Astron. Astrophys. 2004. 42:79–118
doi: 10.1146/annurev.astro.42.053102.134034
c 2004 by Annual Reviews. All rights reserved
Copyright First published online as a Review in Advance on May 19, 2004

THE FIRST STARS
Volker Bromm
Department of Astronomy, Harvard University, Cambridge, Massachusetts 02138;
email: vbromm@cfa.harvard.edu

Richard B. Larson
Department of Astronomy, Yale University, New Haven, Connecticut 06520-8101;
email: larson@astro.yale.edu

Key Words cosmology, intergalactic medium, initial mass function,
star formation, Population III■ Abstract We review recent theoretical results on the formation of the first stars
in the universe, and emphasize related open questions. In particular, we discuss theinitial conditions for Population III star formation, as given by variants of the cold
dark matter cosmology. Numerical simulations have investigated the collapse and thefragmentation of metal-free gas, showing that the first stars were predominantly very
massive. The exact determination of the stellar masses, and the precise form of theprimordial initial mass function, is still hampered by our limited understanding of theaccretion physics and the protostellar feedback effects. We address the importance ofheavy elements in bringing about the transition from an early star formation mode
dominated by massive stars to the familiar mode dominated by low-mass stars at latertimes. We show how complementary observations, both at high redshifts and in our
local cosmic neighborhood, can be utilized to probe the first epoch of star formation.

1. INTRODUCTION
The first stars to form out of unenriched H/He gas marked the crucial transition
from a homogeneous, simple universe to a highly structured, complex one at
the end of the cosmic dark ages. Extending the familiar scheme of classifying
stellar populations in the local universe (Baade 1944) to the extreme case of zerometallicity, the first stars constitute the so-called Population III (e.g., Bond 1981;
Cayrel 1986; Carr 1987, 1994). The quest for Population III stars has fascinatedastronomers for many decades, back to the first tentative ideas of Schwarzschild& Spitzer (1953). Recently, the subject has attracted increased interest, both froma theoretical and an observational perspective. New empirical probes of the highredshift universe have become available, and our ability to carry out sophisticatednumerical simulations has improved dramatically. In this review, we summarize
the current state of this rapidly evolving field.
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The first generation of stars had important effects on subsequent galaxy formation (see Carr et al. 1984 for an early overview). On the one hand, Population
III stars produced copious amounts of UV photons to reionize the universe (e.g.,
Tumlinson & Shull 2000, Bromm et al. 2001b; Schaerer 2002, 2003; Tumlinson
et al. 2003; Venkatesan & Truran 2003; Venkatesan et al. 2003). Recently, theWilkinson Microwave Anisotropy Probe (WMAP) observed the large-angle polarization anisotropy of the cosmic microwave background (CMB), thus constraining
the total ionizing photon production from the first stars (e.g., Cen 2003a,b; Ciardi
et al. 2003; Haiman & Holder 2003; Holder et al. 2003; Kaplinghat et al. 2003;
Kogut et al. 2003; Sokasian et al. 2003a,b; Wyithe & Loeb 2003a,b). The large
optical depth to Thomson scattering measured by WMAP has been interpreted asa signature of a substantial early activity of massive star formation at redshiftsz  15 (see Section 5.1). On the other hand, the supernova (SN) explosions that
ended the lives of the first stars were responsible for the initial enrichment of theintergalactic medium (IGM) with heavy elements (e.g., Ostriker & Gnedin 1996,
Gnedin & Ostriker 1997, Ferrara et al. 2000, Madau et al. 2001, Mori et al. 2002,
Thacker et al. 2002, Bromm et al. 2003, Furlanetto & Loeb 2003, Mackey et al.
2003, Scannapieco et al. 2003, Wada & Venkatesan 2003, Yoshida et al. 2004). An
intriguing possibility unique to zero-metallicity massive stars is the complete disruption of the progenitor in a pair-instability supernova (PISN), which is predicted
to leave no remnant behind (e.g., Barkat et al. 1967, Ober et al. 1983, Bond et al.
1984, Fryer et al. 2001, Heger & Woosley 2002, Heger et al. 2003). This peculiar
explosion mode could have played an important role in quickly seeding the IGM
with the first metals (see Section 4.2).
To place the study of the first stars into the appropriate cosmological context,
one has to ask (see Section 2), When and how did the cosmic dark ages end? During
this time between the emission of the CMB half a million years after the big bang,when the CMB photons shifted into the infrared, and the formation of the firstsources of light, the universe would have appeared completely dark to a humanobserver (see Barkana & Loeb 2001, Loeb & Barkana 2001, Miralda-Escudé 2003
for comprehensive reviews of this cosmic epoch). Historically, the evocative—and
now widely used—dark age analogy was first introduced into the literature briefly
after the ascent of nonbaryonic, cold dark matter (CDM) scenarios (e.g., Sargent
1986, Rees 1993). In the context of popular CDM models of hierarchical structureformation, the first stars are predicted to have formed in dark matter (DM) halosof mass ∼106 M that collapsed at redshifts z  20–30 (e.g., Tegmark et al. 1997,
Yoshida et al. 2003a). The first quasars, however (e.g., Umemura et al. 1993, Loeb
& Rasio 1994, Eisenstein & Loeb 1995), are likely to have formed in more massive
host systems, at redshifts z  10 (Haiman & Loeb 2001, Bromm & Loeb 2003a),and certainly before z ∼ 6.4, which is the redshift of the most distant quasar known(Fan et al. 2003). Here, we do not discuss quasar formation in depth, and we refer
the reader to the references above.
The most fundamental question about the first stars is how massive they typicallywere (see Section 3). Results from recent numerical simulations of the collapse5 Aug 2004 22:8

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and fragmentation of primordial gas clouds suggest that the first stars were predominantly very massive, with typical masses M∗  100 M (Bromm et al. 1999,
2002; Nakamura & Umemura 2001; Abel et al. 2000, 2002). The simulations havetaught us that regardless of the detailed initial conditions, the primordial gas attains
characteristic values of temperature and density. These, in turn, correspond to acharacteristic fragmentation scale, given approximately by the Jeans mass, andthey are explained by the microphysics of molecular hydrogen (H2 ) cooling (seeSection 3.2). The minihalos that host the formation of the first stars have virial
temperatures of order ∼1000 K, below the threshold for cooling owing to atomic
hydrogen lines (∼104 K). Therefore, cooling and fragmentation inside these minihalos are possible only via rotational-vibrational transitions of H2 , which can be
excited even at these low temperatures (see Section 2 for details).
Despite the progress already made, many important questions remain unanswered. A particularly important one, for example, is: What is the functional formof the primordial initial mass function (IMF)? Constraining the characteristic massscale still leaves undetermined the overall range of stellar masses and the powerlaw slope, which is likely to be a function of mass (e.g., Yoshii & Saio 1986;
Larson 1998; Nakamura & Umemura 2001, 2002; Omukai & Yoshii 2003). In
addition, it is presently not known whether binaries or, more generally, clusters ofzero-metallicity stars can form (see Section 3.2). This question has important implications for the star formation process. If Population III star formation typicallyresulted in a binary or multiple system, much of the progenitor cloud’s angular momentum could go into the orbital motion of the stars in such a system. However, if
an isolated formation mode were to predominate, the classical angular momentum
barrier could be much harder to breach (see Larson 2003).
How was the epoch of Population III star formation eventually terminated (seeSection 4)? Two categories of negative feedback effects are likely to be important, the first being radiative and the second chemical in nature (e.g., Ciardi et al.
2000b, Schneider et al. 2002a, Mackey et al. 2003). The radiative feedback consists in soft UV photons produced by the first stars photodissociating the rather
fragile H2 molecules in the surrounding gas, thus suppressing the corresponding
H2 cooling (e.g., Haiman et al. 1997, 2000; Omukai & Nishi 1999; Ciardi et al.
2000a; Nishi & Tashiro 2000; Glover & Brand 2001). Massive Population III stars
could then no longer form. There is some debate, however, about whether theradiative feedback from the first stars was positive or negative (e.g., Ferrara 1998;
Machacek et al. 2001, 2003; Oh 2001; Ricotti et al. 2001; Cen 2003b). The formation of H2 molecules is catalyzed by free electrons, and any process that increased
their presence would therefore also boost the abundance of H2 . For example, anearly background of X-ray photons, either from accretion onto black holes (e.g.,
Madau et al. 2004) or from the remnants of Population III SNe, could ionize thegas and thus increase the abundance of free electrons (e.g., Haiman et al. 2000,
Venkatesan et al. 2001, Glover & Brand 2003, Machacek et al. 2003). This debate is ongoing, and it is too early for definitive conclusions (see Section 4.1 fordetails).

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The second feedback effect that is important in terminating the epoch of Population III star formation is chemical—it is due to the enrichment of the primordialgas with the heavy elements dispersed by the first SNe (see Section 4.2). Numerical simulations of the fragmentation process have shown that lower mass
stars can only form out of gas that was already preenriched to a level in excess of a critical metallicity estimated to be of order 10−4 –10−3 the solar value(Bromm et al. 2001a; see also Schneider et al. 2002a). Depending on the natureof the first SN explosions, and in particular on how efficiently and widespreadthe metal-enriched ejecta are mixed into the surrounding gas, the cosmic starformation will at some point undergo a fundamental transition from an early,
high-mass (Population III) dominated mode to one dominated by lower mass stars(Population II).
We conclude our review with a discussion of observational probes of the natureof the first stars (Section 5). Predicting the properties of the first stars is important
for the design of upcoming instruments, such as the James Webb Space Telescope1(JWST) or the next generation of large (>10 m) ground-based telescopes. The
hope is that over the coming decade, it will become possible to confront current
theoretical predictions about the properties of the first sources of light with observational data. The increasing volume of data on high-redshift galaxies and quasars
from existing ground- and space-based telescopes signals the emergence of thisnew frontier in cosmology.

2. COSMOLOGICAL CONTEXT
The establishment of the current standard CDM model for cosmological structureformation (see Kirshner 2003, Ostriker & Steinhardt 2003 for recent reviews)
has provided a firm framework for the study of the first stars. Within variantsof the CDM model, where larger structures are assembled hierarchically throughsuccessive mergers of smaller building-blocks, the first stars are predicted to form
in DM minihalos of typical mass ∼106 M at redshifts z ∼ 20−30 (e.g., Couchman& Rees 1986). The virial temperatures in these low-mass halos, Tvir ∝ M 2/3 (1 + z)
(Barkana & Loeb 2001), are below the threshold, ∼104 K, for efficient cooling
owing to atomic hydrogen (e.g., Hutchings et al. 2002, Oh & Haiman 2002). It was
realized early on that cooling in the low-temperature primordial gas had to rely on
molecular hydrogen instead (Saslaw & Zipoy 1967, Peebles & Dicke 1968).
Because the thermodynamic behavior of the primordial gas is primarily controlled by H2 cooling, it is crucial to understand the nonequilibrium chemistryof H2 formation and destruction (e.g., Lepp & Shull 1984, Anninos & Norman
1996, Abel et al. 1997, Galli & Palla 1998, Puy & Signore 1999). Because there
are no dust grains to facilitate their formation (e.g., Hirashita & Ferrara 2002),
molecules have to form in the gas phase. The most important formation channel
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See http://ngst.gsfc.nasa.gov.

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turns out to be the sequence H + e− → H− + γ , followed by H− + H → H2 + e−
(McDowell 1961). The free electrons act as catalysts, and they are present as aresidue from the epoch of recombination (Seager et al. 2000) or result from collisional ionization in accretion shocks during the hierarchical build-up of galaxies(e.g., Mac Low & Shull 1986, Shapiro & Kang 1987). The formation of hydrogenmolecules thus ceases when the free electrons have recombined. An alternative
formation channel relies on the intermediary H+
2 , with free protons as catalysts(e.g., Haiman et al. 1996, Abel et al. 1997). The H− channel, however, dominates
in most circumstances (e.g., Tegmark et al. 1997). Calculations of H2 formation incollapsing top-hat overdensities, idealizing the virialization of DM halos in CDM
cosmogonies, have found a simple approximate relationship between the asymp1.5
totic H2 abundance and virial temperature in the overdensity (or halo): f H2 ∝ Tvir
(Tegmark et al. 1997).
Applying the familiar criterion (Rees & Ostriker 1977, Silk 1977) for the formation of galaxies that the cooling timescale has to be shorter than the dynamicaltimescale, tcool < tdyn , one can derive the minimum halo mass at a given redshiftinside of which the gas is able to cool and eventually form stars (e.g., Tegmark
et al. 1997, Santoro & Thomas 2003). The H2 cooling function has been quite
uncertain, with rates that differ by an order of magnitude over the relevant temperature regime (Hollenbach & McKee 1979, Lepp & Shull 1984). Recent advancesin the quantum-mechanical computation of the collisional excitation process (Hatoms colliding with H2 molecules) have provided a much more reliable determination of the H2 cooling function (see Galli & Palla 1998, and references therein).
Combining the CDM prescription for the assembly of virialized DM halos with
the H2 -driven thermal evolution of the primordial gas, a minimum halo mass of∼106 M is required for collapse redshifts z vir ∼ 20–30. From detailed calculations, one finds that the gas in such a “successful” halo has reached a molecule
fraction in excess of f H2 ∼ 10−4 (e.g., Haiman et al. 1996, Tegmark et al. 1997,
Yoshida et al. 2003a). These systems correspond to 3σ –4σ peaks in the Gaussian random field of primordial density fluctuations. In principle, DM halos thatare sufficiently massive to harbor cold, dense gas clouds could form at higher
redshifts, z vir  40. Such systems, however, would correspond to extremely rare,high-σ peaks in the Gaussian density field (e.g., Miralda-Escudé 2003).
To more realistically assess the formation of cold and dense star-forming clouds
in the earliest stages of cosmological structure formation, three-dimensional simulations of the combined evolution of the DM and gas are required within a cosmological set-up (Ostriker & Gnedin 1996, Gnedin & Ostriker 1997, Abel et al.
1998). These studies confirmed the important role of H2 cooling in low-mass halos
at high z. Recently, the problem of forming primordial gas clouds within a fully
cosmological context has been revisited with high numerical resolution (Yoshida
et al. 2003a,b,c,d). The resulting gas density field is shown in Figure 1 for a standard CDM cosmology at z = 17. The bright knots at the intersections of thefilamentary network are the star-forming clouds. They have individual masses (DMand gas) of ∼106 M . As expected from the statistics of the high-σ peaks (e.g.,5 Aug 2004 22:8

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Figure 2 Condition for successful formation of cold, dense gas clouds. Massweighted mean H2 fraction versus virial temperature. (Open circles) DM halos thatfail to host cold gas clouds. (Filled circles) DM halos that succeed in harboring starforming clouds. The critical line corresponds to the condition tcool ∼ tff . (From Yoshida
et al. 2003a.)

Kaiser 1984), the primordial clouds are predominantly clustered, although there
are a few cases of more isolated ones.
Whether a given DM halo successfully hosts a cold, dense (T  0.5 Tvir , n H 
5×102 cm−3 ) gas cloud can nicely be understood with the Rees-Ostriker criterion,
as shown in Figure 2. Yoshida et al. (2003a) derive a minimum collapse mass ofMcrit  7 × 105 M , with only a weak dependence on collapse redshift (seealso Haiman et al. 1996, Fuller & Couchman 2000, Machacek et al. 2001). The
dynamical heating accompanying the merging of DM halos has an important effecton the thermal and chemical evolution of the gas. Clouds do not successfully cool
if they experience too rapid a growth in mass.
The primordial gas clouds that are found in the cosmological simulations are thesites where the first stars form. It is, therefore, important to learn what properties
these clouds have in terms of overall size, shape, and angular momentum contents.
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where L, E, and M are the total angular momentum, energy, and mass, respectively.
The spin parameter is a measure of the degree of rotational support, such that theratio of centrifugal to gravitational acceleration is given by ∼λ2 at virialization.
The spin values measured in pure DM cosmological simulations can be describedby a lognormal distribution function with a mean value of λ̄ = 0.04, similar to what
is found for larger-scale systems (Jang-Condell & Hernquist 2001). Interestingly,
the resulting angular momentum vectors for the gaseous and DM components
generally do not align in the simulations of Yoshida et al. (2003a). The overall
sizes of the Population III star-forming clouds are close to the virial radius of
the host DM halo, with Rvir ∼ 100 pc, which is not too different from the typicaldimensions of present-day giant molecular clouds (e.g., Larson 2003). Dependingon the degree of spin, the clouds have shapes with various degrees of flattening(e.g., Bromm et al. 2002, Yoshida et al. 2003a). Owing to the importance ofpressure forces, however, overall cloud shapes tend to be rather spherical. To
fully elucidate the properties of the star-forming primordial clouds, even higher
resolution cosmological simulations will be necessary.
The theoretical predictions for the formation sites of the first stars sensitively
depend on the exact nature of the DM component and its fluctuation spectrum.
Recently, two models have been discussed that would much reduce the fluctuation
power on small mass scales. The first of these, the warm dark matter (WDM)
model (e.g., Bode et al. 2001), has been proposed to remedy the well-known
problems of standard CDM on subgalactic scales (e.g., Flores & Primack 1994,
Moore et al. 1999). These concern the predicted large abundance of substructureor, equivalently, of satellite systems, and the high (cuspy) densities in the centers ofgalaxies. Both predictions are in conflict with observations. The second model, the“running” spectral index (RSI) model, is suggested by the combined analysis of
the WMAP data, the 2dF galaxy redshift survey, and Lyman-α forest observations(Spergel et al. 2003, Peiris et al. 2003). A series of recent studies have worked
out the consequences of these reduced small-scale power models on early starformation (Somerville et al. 2003; Yoshida et al. 2003b,c). Within these models,
the star formation rate at z  15 is significantly reduced compared to the standard
CDM case. This is due to the absence of low-mass halos and their associated gas
clouds, which are cooled by molecular hydrogen. We revisit this issue when we
discuss the implications of the WMAP data for early star formation (Section 5.1).

3. FORMATION OF THE FIRST STARS
3.1. Star Formation Then and Now
Currently, we do not have direct observational constraints on how the first stars
formed at the end of the cosmic dark ages. It is, therefore, instructive to briefly
summarize what we have learned about star formation in the present-day universe,
where theoretical reasoning is guided by a wealth of observational data (see Pudritz
2002, Ward-Thompson 2002, Larson 2003 for recent reviews).

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Population I stars form out of cold, dense molecular gas that is structured in acomplex, highly inhomogeneous way. The molecular clouds are supported against
gravity by turbulent velocity fields and pervaded on large scales by magnetic
fields. Stars tend to form in clusters ranging from a few hundred up to ∼106 stars.
It appears likely that the clustered nature of star formation leads to complicateddynamics and tidal interactions that transport angular momentum, thus allowing
the collapsing gas to overcome the classical centrifugal barrier (Larson 2002). The
IMF of Population I stars is observed to have the approximate Salpeter (1955)
form (e.g., Kroupa 2002)

where


x

dN
∝ M −x ,d log M(1)

1.35 for M  0.5 M
.
0.0 for M  0.5 M

(2)

The IMF is not observationally well determined at the lowest masses, but theory
predicts that there should be a lower mass limit of approximately 0.007 M set
by opacity effects. This theoretical limit reflects the minimum fragment mass, set
when the rate at which gravitational energy is released during the collapse exceeds
the rate at which the gas can cool (e.g., Low & Lynden-Bell 1976, Rees 1976).
The most important feature of the observed IMF is that ∼1 M is the characteristic
mass scale of Population I star formation in the sense that most of the mass goes
into stars with masses close to this value. Recent hydrodynamical simulations of
the collapse and fragmentation of turbulent molecular cloud cores (e.g., Padoan &Nordlund 2002, Bate et al. 2003) illustrate the highly dynamic and chaotic natureof the star formation process.2
The metal-rich chemistry, magnetohydrodynamics, and radiative transfer involved in present-day star formation is complex, and we still lack a comprehensive
theoretical framework that predicts the IMF from first principles. Star formation in
the high-redshift universe, however, poses a theoretically more tractable problem
owing to a number of simplifying features, such as (a) the initial absence of heavy
elements and therefore of dust; (b) the absence of dynamically significant magnetic
fields in the pristine gas left over from the big bang; and (c) the absence of any
effects from previous episodes of star formation, which would completely alter
the conditions for subsequent star formation. The cooling of the primordial gas
then depends only on hydrogen in its atomic and molecular form. Whereas in thepresent-day interstellar medium, the initial state of the star-forming cloud is poorly
constrained, the corresponding initial conditions for primordial star formation are
simple, given by the popular CDM model of cosmological structure formation.
We now turn to a discussion of this theoretically attractive and important problem.
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3.2. Clump Formation: The Characteristic Stellar Mass
How did the first stars form? This subject has a long and venerable history (e.g.,
Schwarzschild & Spitzer 1953; Matsuda et al. 1969; Yoneyama 1972; Hutchins
1976; Silk 1977, 1983; Yoshii & Sabano 1979; Carlberg 1981; Kashlinsky & Rees
1983; Palla et al. 1983; Yoshii & Saio 1986). In this review, we focus mainly on
the more recent work since the renewed interest in high-redshift star formationthat began in the mid-1990s (e.g., Haiman et al. 1996, Uehara et al. 1996, Haiman& Loeb 1997, Tegmark et al. 1997, Larson 1998).
The complete answer to this question would entail a theoretical prediction for
the Population III IMF, which is rather challenging. A more tractable task is toestimate the characteristic mass scale, Mc , of the first stars; most of the recent
numerical work has focused on this simpler problem. The characteristic mass isthe mass below which the IMF flattens or begins to decline (see Larson 1998
for examples of possible analytic forms). As mentioned above, this mass scale isobserved to be ∼1 M in the present-day universe. Because the detailed shape of
the primordial IMF is highly uncertain, we focus here on constraining Mc , as thismass scale indicates the typical outcome of the primordial star formation process,or where most of the available mass ends up (e.g., Clarke & Bromm 2003).
To fully explore the dynamical, thermal, and chemical properties of primordialgas, three-dimensional numerical simulations are needed, although more idealized
investigations in two, one, or even zero (one-zone models) dimension are importantin that they allow researchers to probe a larger parameter space. The proper initialconditions for primordial star formation are given by the underlying model ofcosmological structure formation (see Section 2). It is, therefore, necessary tosimulate both the DM and gaseous (baryonic) components. To date, two different
numerical approaches have been used to simulate the general three-dimensional
fragmentation problem in its cosmological context.
The first series of simulations used the adaptive mesh refinement (AMR) technique (Abel et al. 2000, 2002). The AMR algorithm creates a nested hierarchy ofever finer grids in fluid regions where high resolution is needed (see, e.g., Balsara
2001 for a recent review). This method allows one to accurately simulate problems
with high dynamic range, i.e., large contrasts in density. Such a situation indeed
occurs in primordial star formation, where the proper initial conditions are givenby the large-scale (  10 kpc) cosmological environment, and where the collapsinggas has to be followed all the way to protostellar scales (  10−3 pc). The second
series of simulations used the smoothed particle hydrodynamics (SPH) method(Bromm et al. 1999, 2002). This Lagrangian technique samples the continuous
fluid with a finite number of particles whose mass is smoothed out according to agiven distribution function (the SPH kernel). It has been applied to a wide range ofproblems in present-day star formation (see Monaghan 1992 for a standard review).
The SPH approach has the important advantage that it can easily accommodate
the creation of sink particles (e.g., Bate et al. 1995). Quite generically, gravitational
collapse proceeds very nonuniformly, such that the density in a few subregions

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becomes much larger than the average. Numerically, this situation necessitates
the adoption of increasingly small timesteps. Although the simulation follows theevolution of the select density peak in detail, the overall system is essentiallyfrozen. To study the collective properties of star formation, i.e., the formation ofmultiple high-density clumps that interact with each other in a complex fashion(see Bate et al. 2003), the SPH particles are merged into more massive ones,
provided they exceed a given density threshold. Recently, the dynamical rangeof the standard SPH method has been significantly improved by implementing
a “particle splitting” technique (Kitsionas & Whitworth 2002, Bromm & Loeb
2003a). When the simulation reaches such high density in a certain region that asink particle would normally be created, a complementary strategy is now adopted:
Every SPH particle in the unrefined, high-density region acts as a parent particle
and spawns a given number of child particles, endowing them with its properties.
In effect, this is the SPH equivalent of the grid-based AMR technique.
The most important difference between these two studies lies in the way the
initial conditions are implemented. The Abel et al. (2000, 2002) simulations startat z = 100 with a realistic cosmological set-up, considering a periodic volumeof physical size 128/(1 + z) kpc. The AMR technique allows Abel et al. (2000,
2002) to bridge the gap from cosmological to protostellar scales. The Bromm
et al. (1999, 2002) effort, however, initializes the simulations, also at z = 100, by
realizing spherical overdensities that correspond to high-σ peaks in the Gaussian
random field of cosmological density fluctuations (see Katz 1991).
The systems considered by Bromm et al. (1999, 2002) comprise halos of total
mass 105 –107 M that collapse at z vir  20–30. The virialization (or collapse)
redshift, z vir , marks the approximate instant where the DM component settles into
virial equilibrium (leading to −E grav ∼ 2E kin ). The Bromm et al. (1999, 2002) initial configurations are not simple “top-hat” (i.e., uniform density) clouds, however,as the dynamically dominant DM component is endowed with density fluctuations
according to P(k) ∝ k −3 , the asymptotic behavior of the CDM power spectrum on small mass scales (e.g., Yoshida et al. 2003d). The Bromm et al. (1999,
2002) simulations, therefore, do reflect the basic bottom-up, hierarchical merging
of smaller DM halos into larger ones, and the accompanying shock heating of thegas. Because of the isolated nature of their systems, Bromm et al. (1999, 2002)
had to treat the initial amount of angular momentum (or spin) as a free parameter,
imparting both the DM and gaseous components with solid-body rotation aroundthe center of mass. In the Abel et al. (2000, 2002) simulations, tidal torques fromneighboring halos generate the halo spin self-consistently. In summary, the Abelet al. (2000, 2002) investigation is the most realistic treatment of the primordial starformation problem to date, whereas the more idealized Bromm et al. (1999, 2002)
simulations, implementing the same physical ingredients as Abel et al. (2000,
2002), have explored a much wider range of initial conditions.
In comparing the simulations of Abel et al. (2000, 2002) and Bromm et al.
(1999, 2002), the most important aspect is that both studies, employing very
different methods, agree on the existence of a preferred state for the primordial gas,corresponding to characteristic values of temperature and density, Tc ∼ 200 K and5 Aug 2004 22:8

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P1: IKH89n c ∼ 104 cm−3 , respectively. These characteristic scales turn out to be rather robustin the sense that they are not very sensitive to variations in the initial conditions.
In Figure 3 (panel c), this preferred state in the T − n phase diagram can clearly bediscerned. This figure, from the simulations of Bromm et al. (1999, 2002), plots
the respective gas properties for each individual SPH particle. The diagram thus
contains an additional dimension of information: Where evolutionary timescales
are short, only a few particles are plotted, whereas they tend to accumulate where
the overall evolution is slow. Such a “loitering” state is reached at Tc and n c .

Figure 3 Properties of primordial gas. (a) Free electron abundance versus hydrogen
number density (in cm−3 ). (b) Hydrogen molecule abundance versus number density.
(c) Gas temperature versus number density. At densities below ∼1 cm−3 , the gas temperature rises because of adiabatic compression until it reaches the virial value of Tvir 5000 K. At higher densities, cooling owing to H2 drives the temperature down again
until the gas settles into a quasi-hydrostatic state at Tc ∼ 200 K and n c ∼ 104 cm−3 .
Upon further compression owing to the onset of the gravitational instability, the temperature experiences a modest rise again. (d) Jeans mass (in M ) versus number density.
The Jeans mass reaches a value of M J ∼ 103 M for the quasi-hydrostatic gas in thecenter of the DM potential well. (From Bromm et al. 2002.)

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These characteristic scales can be understood by considering the microphysicsof H2 , the main coolant in metal-free, star-forming gas (Abel et al. 2002, Bromm
et al. 2002). At temperatures  1000 K, cooling is due to the collisional excitationand subsequent radiative decay of rotational transitions. The two lowest-lying
rotational energy levels in H2 have an energy spacing of E/kB  512 K. Collisions
with particles (mostly H atoms) that populate the high-energy tail of the MaxwellBoltzmann velocity distribution can lead to somewhat lower temperatures, but H2
cooling cannot proceed to T  100 K. This explains the characteristic temperature.
The characteristic density, in turn, is given by the critical density above whichcollisional deexcitations, which do not cool the gas, compete with radiative decays,
which lead to cooling. This saturation of the H2 cooling marks the transition from
non-local thermodynamic equilibrium (non-LTE) rotational level populations tothermal (LTE) ones. At densities below n c , the cooling rate is proportional to thedensity squared, whereas at higher densities, the dependence is only linear.
Once the characteristic state is reached, the evolution toward higher density is
temporarily halted owing to the now inefficient cooling, and the gas undergoesa phase of quasi-hydrostatic, slow contraction. To move away from this loitering
regime, enough mass has to accumulate to trigger a gravitational runaway collapse.
This condition is simply M  M J , where the Jeans mass or, almost equivalently,
the Bonnor-Ebert mass3 can be written as (e.g., Clarke & Bromm 2003)
3/2 
−1/2Tn
M J  700 M
.
(3)
200 K
104 cm−3
Primordial star formation, where magnetic fields and turbulence are expected toinitially play no important dynamical role, may be the best case for the applicationof the classical Jeans criterion, which is describing the balance between gravity
and the opposing thermal pressure (Jeans 1902).
A prestellar clump of mass M  M J is the immediate progenitor of a single
star or, in the case of further subfragmentation, a binary or small multiple system.
In galactic star-forming regions, like ρ Ophiuchi, such clumps with masses close
to stellar values have been observed as gravitationally bound clouds that lackthe emission from embedded stellar sources (e.g., Motte et al. 1998). The highdensity clumps are clearly not stars yet. To probe the further fate of a clump, one
first has to follow the collapse to higher densities up to the formation of an opticallythick hydrostatic core (see Section 3.3), and subsequently the accretion from thediffuse envelope onto the central core (see Section 3.4). The parent clump mass,
however, already sets an upper limit for the final stellar mass whose precise value
is determined by the accretion process.
Although Abel et al. (2000, 2002) and Bromm et al. (1999, 2002) agree on
the magnitude of the characteristic fragmentation scale, Abel et al. (2000, 2002)
3

An equilibrium isothermal sphere of finite size with a fixed temperature and boundarypressure is unstable to collapse if its mass exceeds the critical Bonnor-Ebert value (seeLarson 2003).

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have argued that only one star forms per halo, whereas Bromm et al. (1999, 2002)
have simulated cases where multiple clumps form, such that the number of stars
in a minihalo is N∗ ∼ 1–5. Among the cases studied by Bromm et al. (1999,
2002) favoring N∗ > 1 are high-spin runs that lead to disk-like configurations,which subsequently fragment into a number of clumps or more massive halos (withtotal mass M  107 M ). Notice that the case simulated by Abel et al. (2000, 2002)
corresponds to a gas mass of ∼3×104 M . Such a low-mass halo is only marginallyable to cool, and when Bromm et al. (1999, 2002) simulated a system with thismass, they also found that only a single star forms inside the halo. Further work is
required to elucidate whether the multiple-clump formation in Bromm et al. (1999,
2002) also occurs in simulations with realistic cosmological initial conditions.
Whether primordial clusters of metal-free stars are able to form in the firstgalaxies is an important open question. The answer will determine the properinterpretation of the primordial IMF. If Population III star formation were generally
clustered, as is the case in the present-day universe, the IMF would simply describe
the actual distribution of stellar masses in a given cluster. However, if the firststars formed predominantly in isolation, the IMF would have the meaning of aprobability distribution in a random process that results in only one stellar mass
per “draw.” The present-day IMF, however, appears to be shaped largely by the
chaotic interactions between many accreting protostars in a common reservoir ofgas (e.g., Bate et al. 2003). If there were no cluster to begin with, how would the
IMF be built up? Would there still be a self-similar extension at masses M > Mc ?
Progress has to rely on improved simulations in future work, and it is difficult toguess the outcome.
A possible clustered nature of forming Population III stars would also haveimportant consequences for the transport of angular momentum. In fact, a crucial
question is how the primordial gas can so efficiently shed its angular momentum
on the way to forming massive clumps. In the case of a clustered formation process, angular momentum could be efficiently transported outward by gravitational(tidal) torques. Tidal torques can then transfer much of the angular momentum
from the gas around each forming clump to the orbital motion of the system,
similar to the case of present-day star formation in a clustered environment (e.g.,Larson 2002, Bate et al. 2003). In the isolated formation process of Abel et al.
(2000, 2002), however, such a mechanism is not available. Alternatively, Abelet al. (2000, 2002) have suggested the efficient transport of angular momentum via
hydrodynamic shocks during turbulent collapse. Again, more work is required toconvincingly sort out the angular momentum transport issue. It may be possiblethat magnetic fields in an accretion disk around the primordial protostar could be
sufficiently amplified by dynamo action from the suspected very weak primordial
values (e.g. Kulsrud et al. 1997). The dynamo activity could occur as a consequence of turbulence that is driven by gravitational instabilities in the disk (Tan &Blackman 2004), and magnetic stresses could then facilitate angular momentum
transport in the accretion disk.
As pointed out above, important contributions have also been made by more
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the DM component, thus implicitly assuming that the gas has already dissipatively
collapsed to the point where the DM ceases to be dynamically dominant. Much attention has been paid to the collapse of filamentary clouds (Nakamura & Umemura
1999, 2001, 2002). Results from one- and two-dimensional simulations that otherwise include all the relevant processes for the primordial chemistry and cooling
have estimated fragmentation scales Mc ∼ 1 M and Mc ∼ 100 M , dependingon the (central) initial density. When the initial density is larger than ∼105 cm−3 ,fragmentation leads to the lower mass value. This bifurcation has led to the prediction of a bimodal IMF for the first stars (Nakamura & Umemura 2001). It is,however, not obvious how such a high initial density can be reached in a realistic
situation where the collapse starts from densities that are typically much smaller
than the bifurcation value.
Recently, the question of whether low-mass metal-free stars can form has received considerable interest in connection with the discovery of extremely metalpoor (or more precisely iron-poor) galactic halo stars (see Section 5.2). Is it possible
to place a firm lower limit on the mass of the first stars? Such a minimum mass may
be given by the so-called opacity limit for fragmentation (e.g., Low & Lynden-Bell
1976, Rees 1976). The opacity limit, estimated to be M F ∼ 0.007 M in presentday star formation, is the mass of a cloud that is no longer able to radiate away its
gravitational binding energy in a free-fall time. For the case of metal-free stars,
one-dimensional, spherically symmetric calculations have indicated a value thatis rather similar to the present-day opacity limit (Omukai & Nishi 1998; see also
Section 3.3). Whereas Omukai & Nishi (1998) have included all relevant sourcesof opacity in metal-free gas, a somewhat more idealized study by Uehara et al.
(1996) has estimated that a cloud of ∼1 M may become opaque to H2 -line radiation. It is currently not clear how relevant the latter value is in a realistic collapse
calculation.
We conclude this section with three other recent suggestions of how to produce
lower-mass stars out of zero-metallicity gas. The first one is a modified version of
the double-peaked (bimodal) IMF and it relies on the presence of deuterium hydride(HD) (Uehara & Inutsuka 2000, Nakamura & Umemura 2002). In principle, HD
is a much more efficient coolant than H2 . It possesses a permanent electric dipole
moment with Einstein A coefficients that are larger by a factor of 103 comparedto the quadrupole transition probability for H2 . It can cool the gas to temperatures
below ∼100 K because its rotational level spacings are smaller than those of H2(e.g., Puy & Signore 1996, Flower et al. 2000). Bromm et al. (2002) included
HD cooling in their simulations and found that it does not change the thermal
evolution of the gas. It may, however, be possible to reach a regime where HD
cooling becomes important when the gas cools down from temperatures above
10,000 K, e.g., in photoionization heated gas or in shock-heated material in dwarfgalaxy-size DM halos (masses 108 M ).
The second idea postulates that the fragmentation scale (in effect the Jeansmass) will be reduced in gas that is bathed in soft UV radiation below the highmasses predicted for the first stars (Omukai & Yoshii 2003). It is argued that the5 Aug 2004 22:8

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strong UV radiation in the vicinity of a massive Population III star will destroy H2 ,thus preventing the gas from reaching the characteristic loitering state. The gas
may then be able to collapse, albeit at higher temperatures, to increasingly large
densities (see also Omukai 2001). The three-dimensional simulations of Bromm& Loeb (2003a), although treating a very similar case, have failed to show thepostulated fragmentation to small masses.
The final idea proposed the formation of lower-mass stars in zero-metal, shockcompressed gas (Mackey et al. 2003). The required compression could, for example, be induced by the first energetic SN explosions. Assuming that the shock isradiative and isobaric (e.g., Shapiro & Kang 1987) one can derive the expression
−1
−1/2 u shn0
M J  10 M
,(4)
102 cm−3
200 kms−1
where n 0 denotes the preshock density and u sh the velocity of the shock. Indeed,because a metal-free population of stars with a characteristic mass intermediate
between the “classical” Population II and Population III cases would be quitedistinct, Mackey et al. (2003) have suggested naming these hypothetical stars
Population II.5. Numerical simulations of SN explosions at high redshifts have so
far failed to bear out this prediction (Bromm et al. 2003), but it may be possiblethat Population II.5 stars can only form in more massive systems (see Salvaterra
et al. 2003).
We conclude this section by pointing out that all the realistic three-dimensional
simulations of the first star formation problem to date have consistently shownthat the primordial gas fragments into massive clumps. Conversely, none of thesuggestions for making low-mass stars out of zero-metallicity gas have yet beenrealized in a three-dimensional simulation. The tentative conclusion, therefore, isthat the typical outcome of primordial star formation consists in massive clumps,
and that low-mass clumps may need very special conditions to form.

3.3. Protostellar Collapse
What is the further fate of the clumps discussed above? In particular, one would liketo test the notion that such a Jeans-unstable clump is the immediate progenitor of asingle star that forms in its center. This evidently will only be correct if the clump
does not undergo further subfragmentation upon collapsing to higher densities. It
has long been suspected that such subfragmentation could occur at densities inexcess of ∼108 cm−3 , at which point three-body reactions become very efficientin converting the atomic gas (with only a trace amount of H2 molecules from the
H− channel) into fully molecular form: 3H → H2 + H (Palla et al. 1983). The
H2 coolant suddenly gets more abundant by a factor of ∼103 . The corresponding
boost in cooling could trigger a thermal instability, thus breaking up the clump into
smaller pieces (Silk 1983). Both Abel et al. (2000, 2002) and Bromm et al. (1999,
2002) have included the three-body reactions in their chemical reaction networks
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subfragmentation occurs [for the Bromm et al. (1999, 2002) model, this calculation
is described in Bromm 2000 and Bromm & Loeb 2004]. Both groups report that nofurther subfragmentation is seen. With hindsight, this may not be too surprising.
The reason is that any small density fluctuations that are present earlier on andthat could serve as seeds for later fragmentation will have been erased by pressure
forces during the slow, quasi-hydrostatic loitering phase at n ∼ n c . In addition,inefficient cooling may also play a role in suppressing high-density fragmentation.
Despite the increase in the cooling rate throughout the fully molecular gas, coolingnever leads to a significant drop in temperature owing to the counteracting effectof compressional heating and H2 formation heating.
Extending the analogous calculation for the collapse of a present-day protostar(Larson 1969) to the primordial case, Omukai & Nishi (1998) have carried out
one-dimensional hydrodynamical simulations in spherical symmetry. They also
consider the full set of chemical reactions and implement an algorithm to solve
for the radiative transfer in the H2 lines as well as in the continuum. The mostimportant result is that the mass of the initial hydrostatic core, formed in the centerof the collapsing cloud when the density is high enough (n ∼ 1022 cm−3 ) for thegas to become optically thick to continuum radiation, is almost the same as inpresent-day star formation: Mcore ∼ 5 × 10−3 M . In Figure 4, the radial profilesof density, temperature, velocity, and H2 abundance are shown (reproduced from
Omukai & Nishi 1998). The profiles of density and velocity before the time ofcore formation (corresponding to the curves with label 6) are well described by
the Larson-Penston (LP) similarity solution (Larson 1969, Penston 1969). The LP
solution approximately describes the growth of the central density peak and isvalid asymptotically in the limit of high central densities and small radii where theinitial and boundary conditions have been “forgotten.” Once the core has formed,
the self-similarity is broken. Similar results have been found by Ripamonti et al.
(2002) who have in addition worked out the spectrum of the radiation that escapesfrom the collapsing clump (mostly in the IR as continuum and line photons). This
type of approximately self-similar behavior seems to be a very generic result of
collapse with a simple equation of state, even when rotation and magnetic fields
are included (see Larson 2003). The apparent robustness of the LP solution thus
supports the results found for the Population III case.
The small value of the initial core does not mean that this will be the final mass of
a Population III star. This instead will be determined by how efficient the accretion
process will be in incorporating the clump mass into the growing protostar. Before
we discuss the accretion process in the following section, we briefly mention thatfragments with masses close to the opacity limit could survive if the accretion
process were somehow prematurely curtailed. This mechanism has been found in
the simulation of a collapsing present-day, turbulent molecular cloud (Bate et al.
2003). A fraction of the accreting cores is ejected, by slingshot interactions withneighboring ones, from the natal cloud before they could grow to stellar masses,
and thus end up as brown dwarfs (Bate et al. 2002). Again, such a mechanism would
only work in a clustered environment, and as discussed above, this might not bevalid for primordial star formation. An alternative suggestion to produce primordial

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Figure 4 Collapse of a primordial protostar. Shown is the evolutionary sequence as
obtained in a one-dimensional calculation, with time progressing from curves labeled
1 to 7. (a) Number density versus radial distance. (b) Gas temperature versus distance.
(c) Velocity versus distance. (d) Hydrogen molecule abundance versus distance. The
curves labeled 6 correspond to the situation briefly after the formation of a centralhydrostatic core. (From Omukai & Nishi 1998.)

brown dwarfs invokes cooling owing to HD, which, as pointed out above, could
lead to lower temperatures and higher densities (Uehara & Inutsuka 2000).

3.4. Accretion Physics
How massive were the first stars? Star formation typically proceeds from theinside-out through the accretion of gas onto a central hydrostatic core. Whereas
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star formation (see above), the accretion process—ultimately responsible for setting the final stellar mass—is expected to be rather different. On dimensional
grounds, the accretion rate is simply related to the sound speed cubed over Newton’sconstant (or equivalently given by the ratio of the Jeans mass and the free-fall time):Ṁ acc ∼ cs3 /G ∝ T 3/2 . A simple comparison of the temperatures in present-day
star-forming regions (T ∼ 10 K) with those in primordial ones (T ∼ 200–300 K)
already indicates a difference in the accretion rate of more than two orders ofmagnitude.
Extending earlier work by Stahler et al. (1986), Omukai & Palla (2001, 2003)
have investigated the accretion problem in considerable detail, going beyond thesimple dimensional argument given above. Their computational technique approximates the time evolution by considering a sequence of steady-state accretion
flows onto a growing hydrostatic core. Somewhat counterintuitively, these authors identify a critical accretion rate, Ṁ crit ∼ 4 × 10−3 M yr−1 , such that for
accretion rates higher than this, the protostar cannot grow to masses much in excess of a few tens of solar masses. For smaller rates, however, the accretion ispredicted to proceed all the way up to ∼600 M , i.e., on the order of the hostclump.
The physical basis for the critical accretion rate is that for ongoing accretion
onto the core, the luminosity must not exceed the Eddington value, L EDD . In theearly stages of accretion, before the onset of hydrogen burning, the luminosity isapproximately given by L tot ∼ L acc  G M∗ Ṁ acc /R∗ . By demanding L acc  L EDD
it follows
Ṁ crit 

L EDD R∗∼ 5 × 10−3 M year−1 ,
G M∗

(5)

where R∗ ∼ 5R , a typical value for a Population III main-sequence star (e.g.,
Bromm et al. 2001b). In Figure 5 (from Omukai & Palla 2003), the mass-radiusrelation is shown for various values of the accretion rate. The dramatic swellingin radius effectively shuts off accretion at M∗  100 M , when the accretion rate
exceeds Mcrit .
Realistic accretion flows are expected to have a time-dependent rate, and theoutcome will thus depend on whether the accretion rate will decline rapidly enoughto avoid exceeding the Eddington luminosity at some stage during the evolution.
The biggest caveat concerning the Omukai & Palla (2003) results seems to be theissue of geometry. A three-dimensional accretion flow of gas with some residual
degree of angular momentum will deviate from spherical symmetry and instead
form a disk. It is then conceivable that most of the photons can escape along theaxes, whereas mass can flow in unimpeded through the accretion disk (see Tan &McKee 2004).
Recently, Bromm & Loeb (2003b, 2004) have used the SPH particle-splitting
technique to study the three-dimensional accretion flow around a primordial protostar. Whereas Bromm et al. (1999, 2002) created sink particles already atn  108 cm−3 , briefly before three-body reactions would convert the gas into

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Figure 5 Evolution of accreting metal-free protostar. Shown is the radius-mass relation for different values of the accretion rate (increasing from bottom to top). Accretionis effectively shut off for the cases with Ṁ  Ṁ crit because of the dramatic increase inradius. (From Omukai & Palla 2003.)

fully molecular form, the gas is now allowed to reach densities of 1012 cm−3 before being incorporated into a central sink particle. Bromm & Loeb (2004) studied
how the molecular core grows in mass over the first ∼104 years after its formation.
The accretion rate is initially very high, Ṁ acc ∼ 0.1 M year−1 , and subsequently
declines in time in a complex fasion. Expressed in terms of the sound speed inPopulation III prestellar cores, this initial rate is Ṁ acc ∼ 25cs3 /G. Therefore, as
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simulations of present-day star formation, the Population III case might again be
an example of a very generic type of behavior for collapse and accretion (seeLarson 2003).
The mass of the molecular core, taken as an estimator of the protostellar
mass, grows approximately as M∗ ∼ Ṁ acc dt ∝ t 0.75 , for t < 103 years, and asM∗ ∝ t 0.4 afterward. The early time accretion behavior is similar to the power-law
scalings found in one-dimensional, spherically symmetric simulations (Omukai &Nishi 1998, Ripamonti et al. 2002). The complex accretion history encountered in
the simulation of Bromm & Loeb (2004), however, cannot be accommodated witha single power-law. Such a multiple power-law behavior has also been inferredin other studies (Abel et al. 2002, Omukai & Palla 2003). This earlier inference,
however, was indirect in that it did not follow the actual accretion flow but instead
estimated the accretion rate from the instantaneous density and velocity profiles
at the termination of the three-dimensional simulation of Abel et al. (2002). In thesimulation of Bromm & Loeb (2004), the accretion rate quickly drops, after only
∼3000 years, below Ṁ crit , at which point the stellar mass is M∗ ∼ 30 M . The
accretion rate thus becomes subcritical early on, at which time the protostar still
undergoes Kelvin-Helmholtz contraction (the descending portion of the curves in
Figure 5). The dramatic expansion in radius, which according to Omukai & Palla(2003) would shut off further accretion, could therefore be avoided. A conservative
upper limit for the final mass of the star is then M∗ (t = 3 × 106 years) ∼ 500 M .
This upper bound has been derived by assuming that accretion cannot go on for
longer than the total lifetime of a very massive star (VMS), which is almost independent of stellar mass (e.g., Bond et al. 1984).
Can a Population III star ever reach this asymptotic mass limit? The answer
to this question is not yet known with any certainty, and it depends on whether
the accretion from the dust-free envelope is eventually terminated by feedback
from the star (e.g., Omukai & Palla 2001, 2003; Ripamonti et al. 2002; Omukai &Inutsuka 2002; Tan & McKee 2004). The standard mechanism by which accretion
may be terminated in metal-rich gas, namely radiation pressure on dust grains(Wolfire & Cassinelli 1987), is evidently not effective for gas with a primordial
composition. Recently, it has been speculated that accretion could instead be turned
off through the formation of an H II region (Omukai & Inutsuka 2002) or throughthe radiation pressure exerted by trapped Lyα photons (Tan & McKee 2004). The
termination of the accretion process defines the current unsolved frontier in studies
of Population III star formation.

3.5. The Second Generation of Stars: Critical Metallicity
How and when did the transition take place from the early formation of massive
stars to that of low-mass stars at later times? The very first stars, marking the cosmic Renaissance of structure formation, formed under conditions that were much
simpler than the highly complex environment in present-day molecular clouds.
Subsequently, however, the situation rapidly became more complicated, again due

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to the feedback from the first stars on the IGM, owing to the production of photons
and of heavy elements (see Section 4).
In contrast to the formation mode of massive stars (Population III) at high redshifts, fragmentation is observed to favor stars below a solar mass (Population
I and II) in the present-day universe. The transition between these fundamental
modes is mainly expected to be driven by the progressive enrichment of the cosmic
gas with heavy elements (see Section 4.2). This enables the gas to cool to lower
temperatures. The concept of a critical metallicity, Z crit , has been used to characterize the transition between Population III and Population II formation modes,
where Z denotes the mass fraction contributed by all heavy elements (Omukai
2000; Bromm et al. 2001a; Schneider et al. 2002a, 2003a; Mackey et al. 2003).
Published studies have only constrained this important parameter to within a few
orders of magnitude, Z crit ∼ 10−6 –10−3 Z  , under the implicit assumption of solar relative abundances of metals. This assumption is likely to be violated by themetal yields of the first SNe at high redshifts, for which strong deviations from solar
abundance ratios are predicted (e.g., Heger & Woosley 2002; Qian & Wasserburg
2002; Qian et al. 2002; Umeda & Nomoto 2002, 2003; Sneden & Cowan 2003).
The cooling rate of the metals depends on their ionization state, which is controlled
by the ionizing backgrounds (UV and X-ray photons or cosmic rays) and which isstill not well known (see Section 4.1).
Recently, Bromm & Loeb (2003c) have shown that the transition between theabove star formation modes is driven primarily by fine-structure line cooling of
singly ionized carbon or neutral atomic oxygen. Earlier estimates of Z crit , which
did not explicitly distinguish between different coolants, are refined by introducing
separate critical abundances for carbon and oxygen, [C/H]crit and [O/H]crit , respectively, where [A/H]= log10 (NA /NH ) − log10 (NA /NH ) , and a subscript  denotes
solar values. Because C and O are also the most important coolants throughout
most of the cool atomic interstellar medium (ISM) in present-day galaxies, it is
not implausible that these species might be responsible for the global shift inthe star formation mode. Indeed, under the temperature and density conditions
that characterize Population III star formation, the most important coolants are
O I and C II, whose fine-structure lines dominate over all other metal transitions(Hollenbach & McKee 1989). Cooling owing to molecules becomes important
only at lower temperatures, and cooling owing to dust grains only at higher densities (e.g., Omukai 2000, Schneider et al. 2003a). The presence of dust is likely tomodify the equation of state at these high densities in important ways (Schneider
et al. 2003a), and it will be interesting to explore its role in future collapse calculations. The physical nature of the dust, however, that is produced by the first SNe iscurrently still rather uncertain (e.g., Loeb & Haiman 1997, Todini & Ferrara 2001,
Nozawa et al. 2003, Schneider et al. 2003b).
A strong UV flux just below the Lyman limit (hν < 13.6 eV) is predicted tobe emitted by the same stars that are responsible for the reionization of the IGM(see Section 4.1). This soft UV radiation can penetrate the neutral IGM and ionize
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11.26 eV. Because the abundance of carbon is so low, the UV background canionize carbon throughout the universe well before hydrogen reionization. Oxygen,
however, will predominantly be neutral prior to reionization because its ionization
potential is 13.62 eV. Cooling is mediated by excitations owing to collisions of
the respective metal with free electrons or hydrogen atoms. At the low fractional
abundances of electrons, xe  10−4 , expected in the neutral (rapidly recombining)gas at z ∼ 15, collisions with hydrogen atoms dominate. This renders the analysis
independent of the very uncertain hydrogen-ionizing backgrounds at high redshifts(cosmic rays or soft X-rays).
To derive the critical C and O abundances, Bromm & Loeb (2003c) start withthe characteristic state reached in primordial gas that cools only through molecularhydrogen (Section 3.2). For the gas to fragment further, additional cooling owing
to C II or O I is required. Fragmentation requires that the radiative cooling rate behigher than the free-fall compressional heating rate. The critical metal abundances
can thus be found by equating the two rates: C II,O I (n, T )  1.5nkB T /tff , where
kB is Boltzmann’s constant. The critical C and O abundances are found by settingn ∼ n c and T ∼ Tc , resulting in [C/H]crit  −3.5 ± 0.1 and [O/H]crit  −3.05 ±0.2. Strictly speaking, these threshold levels are the required abundances in thegas phase.
Even if sufficient C or O atoms are present to further cool the gas, there
will be a minimum attainable temperature that is set by the interaction of theatoms with the thermal CMB: TCMB = 2.7 K(1 + z) (e.g., Larson 1998, Clarke& Bromm 2003). At z  15, this results in a characteristic stellar mass ofM∗ ∼ 20 M (n f /104 cm−3 )−1/2 , where n f > 104 cm−3 is the density at which
opacity prevents further fragmentation (e.g., Rees 1976). It is possible that thetransition from the high-mass to the low-mass star formation mode was modulatedby the CMB temperature and was therefore gradual, involving intermediate-mass(Population II.5) stars at intermediate redshifts (Mackey et al. 2003). This transitional population could give rise to the faint SNe that have been proposed to
explain the observed abundance patterns in metal-poor stars (Umeda & Nomoto
2002, 2003).
When and how uniformly the transition in the cosmic star formation mode
took place was governed by the detailed enrichment history of the IGM. This, inturn, was determined by the hydrodynamical transport and mixing of metals fromthe first SN explosions. We discuss the SN-driven dispersal of heavy elements in
Section 4.2.

4. FEEDBACK EFFECTS ON THE IGM
4.1. Radiative Feedback: The Suicidal Nature of the First Stars
The radiation produced by Population III star formation will influence the subsequent thermal history of the IGM and will modify the properties of star-forminggas. This radiative feedback may occur in a variety of ways that depend on the5 Aug 2004 22:8

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energy range of the stellar photons. Molecular hydrogen is the main coolant in lowtemperature, metal-free gas. It enables the formation of the first stars at z  20–30.
We therefore first focus on the radiative effects that are able to influence the H2chemistry. This feedback could either be negative, if it is due to soft UV photons,
or positive, if it is due to X-rays.
We begin by discussing negative feedback. Molecular hydrogen is fragile andcan readily be destroyed by photons in the Lyman-Werner (LW) bands, within theenergy range 11.2–13.6 eV, via the two-step Solomon process (Stecher & Williams
1967)
H2 + γ → H∗2 → 2H.

(6)

The intermediate stage involves an excited electronic state, H∗2 , from which afraction of the subsequent decays end in the vibrational continuum of the groundstate, resulting in the dissociation of the molecule (e.g., Glover & Brand 2001).
The question arises whether H2 cooling can be suppressed in halos that virialize
after the first stars have formed and built up a background of LW photons (e.g.,
Haiman et al. 1997). These photons have energies just below the Lyman limit
and so can easily penetrate a still-neutral IGM prior to reionization. These secondgeneration halos are typically more massive compared to the ∼105 –106 M systems
that host the formation of the first stars at z  20. The gas might then be able toself-shield against the photo-dissociating LW background (e.g., Glover & Brand
2001, Kitayama et al. 2001, Machacek et al. 2001). The effect of self-shielding isoften taken into account, both in simulations and in analytical work, by writing
the H2 photodissociation rate as kdiss ∝ J21 f shield , with a proportionality constantthat depends on the spectrum of the background radiation. Here, J21 measures theflux at the Lyman limit in the customary units of 10−21 ergs s−1 cm−2 Hz−1 sr−1 .
For the shielding factor, the following approximate expression is usually used:f shield  min[1, (NH2 /1014 cm−2 )−0.75 ] (Draine & Bertoldi 1996). This formula
is accurate only for a static medium, and it will grossly overestimate the H2 line
opacity in the presence of large-scale velocity gradients. In cases with strong bulk
flows, the effect of self-shielding is conservatively overestimated.
The overall normalization of the LW background and its evolution with redshift
are still rather uncertain. Close to the epoch of reionization, however, a significant
flux in the LW bands (hν < 13.6 eV) is expected (e.g., Bromm & Loeb 2003a). The−
flux in the LW bands just below the Lyman limit, J21, could be much larger than+that just above, J21 . Assuming that only a fraction, f esc , of the ionizing photons can−+escape from the star-forming halos, one has J21
 J21/ f esc . This could be quite
large because f esc is expected to be low at high redshifts (Wood & Loeb 2000,
and references therein). Alternatively, a strong background of photo-dissociating
photons could also arise internally from star formation in an early dwarf galaxy
itself (e.g., Omukai & Nishi 1999, Oh & Haiman 2002).
The strong UV background flux estimated above makes it possible for H2
formation to be effectively suppressed close to the redshift of reionization. In theabsence of molecular hydrogen, however, cooling can still proceed via atomic

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transitions in halos of mass (see Barkana & Loeb 2001)

1 + z −3/2
M  108 M
.
10

(7)

The virial temperature, Tvir ∼ 104 K, in these more massive halos allows for thevery efficient cooling of the gas via lines of atomic hydrogen (e.g., Madau et al.
2001, Bromm & Clarke 2002). In cases where the gas temperature is close to thevirial temperature, the gas cloud as a whole can undergo collapse but it will not beable to fragment until high-enough densities are reached so that the Jeans mass has
declined sufficiently (Oh & Haiman 2002, Bromm & Loeb 2003a). Recent numerical simulations of the collapse of a dwarf galaxy at z  10, massive enough foratomic H cooling to operate but where cooling owing to H2 is effectively suppressedby a strong LW background, have shown that very massive (∼106 M ), compact(radii 1 pc) gas clouds form in the center of the DM potential well (Bromm &Loeb 2003a). These clouds may be the immediate progenitors of supermassive
black holes that could seed high-redshift quasar activity.
In contrast to this negative feedback on H2 , opposing positive feedbacks havebeen proposed. As discussed in Section 2, H2 formation is primarily facilitated byfree electrons. Any process that temporarily enhances their abundance, therefore,
will also tend to increase the H2 fraction. Various mechanisms have been suggested. First, X-ray photons from SN remnants (e.g., Oh 2001, Cen 2003a) or theaccretion onto black holes could ionize hydrogen in dense regions, enabling thereformation of H2 (Glover & Brand 2003, Machacek et al. 2003). Second, the freeelectrons in relic H II regions (Oh & Haiman 2003), or inside the outer boundary ofa still-active one (Ricotti et al. 2001, 2002a,b), may accomplish the same. Finally,
the collisional ionization in shock waves, and the subsequent nonequilibrium cooling and molecule reformation has been discussed (e.g., Mac Low & Shull 1986,
Shapiro & Kang 1987, Ferrara 1998, Bromm et al. 2003). The last mechanism
evidently is not strictly radiative in nature, but it seems appropriate to include it
here. In summary, conclusions on the overall sign of the radiative feedback on the
H2 chemistry are still very uncertain. To reach more robust results, sophisticatedcosmological simulations are required, treating the build-up of the radiation backgrounds and their influence on the overall star formation rate in a self-consistentmanner. Such challenging simulations are currently still beyond our capabilities.
A qualitatively different feedback is photoheating owing to radiation with energies high enough to ionize hydrogen atoms (see Barkana & Loeb 2001 for acomprehensive discussion). Because photoionization-heating raises the gas temperature to 104 K, this gas is then evaporated out of the high-redshift potential
wells with virial temperatures Tvir < 104 K. Star formation is thus temporarily shutoff. This photoevaporation feedback could be due to sources in the halo itself (e.g.,
Bromm et al. 2003, Oh & Haiman 2003, Whalen et al. 2003) or to an external ionizing background (e.g., Barkana & Loeb 1999, Susa & Umemura 2000, Kitayama
et al. 2001, Dijkstra et al. 2004, Shapiro et al. 2004, Tassis et al. 2003). External
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to complete reionization; in contrast, internal feedback could operate already atmuch higher redshifts.

4.2. Chemical Feedback: First Supernovae
and Early Metal Enrichment
The crucial transition from a smooth, homogeneous universe to an increasingly
complex and structured one is marked by the death of the first stars. The SN
explosions at the end of the initial burst of Population III star formation disperse thefirst elements into the hitherto pristine IGM. How did this initial metal enrichment
proceed? To account for the widespread presence of metals in the Lyα forest atz ∼ 4–5 (e.g., Songaila 2001, Schaye et al. 2003), star formation in low-masssystems at z  10 has been proposed (e.g., Madau et al. 2001, Mori et al. 2002,
Wada & Venkatesan 2003). The metals produced in these low-mass halos can moreeasily escape from their shallow potential wells than those at lower redshift (e.g.,
Aguirre et al. 2001a,b). In addition, the enriched gas has to travel much shorter
distances between neighboring halos at these early times, and it might therefore
have been easier to establish a uniform metal distribution in the IGM.
The important question arises: How did the first stars die? The answer to thisquestion sensitively depends on the precise mass of a Population III star. In particular, if the star has a mass in the narrow interval 140 M  M∗  260 M , it will
explode as a PISN, leading to the complete disruption of the progenitor (e.g., Fryer
et al. 2001, Heger et al. 2003). Population III stars with masses below or above the
PISN range are predicted to form black holes. This latter fate is not accompanied
by a significant dispersal of heavy elements into the IGM because most of the
newly synthesized metals will be locked up in the black hole. The PISN, however,will contribute all its heavy element production to the surrounding gas.
As we have discussed in Section 3.5, the IGM metal enrichment history is
crucial for understanding the detailed transition in the star formation mode fromhigh-mass to low-mass dominated. If this enrichment is very uniform, the transition
would occur rather suddenly. If, on the other hand, enrichment is very patchy,with islands of high-metallicity gas embedded in mostly still pristine material,
then the transition would occur in a nonsynchronized way, probably occupying aconsiderable time interval (e.g., Scannapieco et al. 2003).
Recently, Bromm et al. (2003) have presented numerical simulations of the firstSN explosions at high redshifts (z ∼ 20). In contrast to earlier work, both analytical (e.g., Larson 1974, Dekel & Silk 1986, Scannapieco et al. 2002, Furlanetto &Loeb 2003) and numerical (e.g., Mori et al. 2002, Thacker et al. 2002, Wada &Venkatesan 2003), their focus is on the minihalos that are the sites for the formationof the very first stars. These halos, with masses of M ∼ 105 –106 M and collapse
redshifts of z  20 are the sites for the formation of the truly first stars. It is not obvious, however, whether star formation in more massive systems would not alreadybe significantly altered by the activity of previous generations of stars. In particular,
PISN explosions may only have occurred in the minihalos mentioned above.

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Bromm et al. (2003) assume that there is one single Population III star in thecenter of the minihalo and that this star is massive enough to explode as a PISN.
To span the plausible energy range, explosion energies of E SN = 1051 and 1053
ergs are considered, corresponding to stars with masses M∗  140 and 260 M ,respectively (Fryer et al. 2001). The latter case marks the upper mass limit of the
PISN regime. The explosion energy is inserted as thermal energy, such that the
explosion is initialized with conditions appropriate for the adiabatic Sedov-Taylor(ST) phase. Because the PISN is completely disrupted without leaving a remnant
behind, it is possible to trace the subsequent fate of the metals by using the SPH
particles that represent the stellar ejecta as markers.
Initially, the blast wave evolves into a roughly uniform medium at radii r Rcore ∼ 20 pc, resulting in the usual ST scaling: Rsh ∝ t 0.4 . Once the blast wave
reaches beyond the core, however, it encounters an isothermal density profile in theremainder of the halo, leading to the scaling (e.g., Ostriker & McKee 1988): Rsh 
30 pc(E SN /1053 ergs)1/3 (t/105 year)2/3 . A few million years after the explosion, radiative (inverse Compton) losses become important, and the SN remnant enters its
final, snowplow phase. The significance of inverse Compton cooling distinguishes
very-high-z SN explosions from those occurring in the present-day universe.
After approximately 105 years, when tcool  Rsh /u sh , a dense shell begins toform at a radius ∼50 pc. The thermodynamic properties of the dense shell, theinterior bubble, and the surrounding medium are summarized in Figure 6. The
shocked, swept-up gas exhibits three distinct thermal phases, as can clearly beseen in Figure 6. One phase is at T  105 K, the second is at T ∼ 104 K, andthe third is at T  103 K. The latter two phases correspond to the dense shell.
The cooling from very high temperatures to ∼104 K and subsequently to 102 K
proceeds in nonequilibrium, such that a fraction of free electrons persists below∼8000 K, allowing the reformation of H2 (e.g., Mac Low & Shull 1986, Shapiro& Kang 1987, Oh & Haiman 2002). The molecular fraction in the densest regions
asymptotically assumes a plateau value of f H2 ∼ 3 × 10−3 . Inside the expanding
SN remnant lies a hot pressurized bubble (the grey symbols in Figure 6) that
corresponds to the original stellar ejecta. This bubble drives the outward motion,and its gas is slowly cooled by adiabatic expansion and inverse Compton losses.
In Figure 7, the projected gas density is shown ∼106 years after the explosion
for explosion energy E SN = 1053 ergs. The blast wave has succeeded in completely
disrupting the minihalo, dispersing most of the halo gas into the IGM. Comparing
E SN with the binding energy of the minihalo, E grav ∼ 1051 ergs, such an outcome
is clearly expected. It depends, however, on the inefficient radiative processes thatcool the hot bubble gas in the initial stages of expansion.
The second prominent feature in Figure 7 is the vigorous fragmentation of
the dense shell into a large number of small cloudlets. The emergence of thesecloudlets, corresponding to gas at T  103 K, is not driven by gravity (e.g.,
Elmegreen 1994, and references therein), but is instead owing to a thermal instability triggered by the onset of atomic cooling. The explosions simulated by Bromm
et al. (2003) do not result in the formation of lower-mass stars, or Population II.5

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Figure 6 Thermodynamic properties of the shocked gas ∼105 years after the explosion of a PISN. Gas temperature (in K) versus log n (in cm−3 ). Dashed line: The virial
temperature of the minihalo. It is evident that Tvir ∼ 3000 K is much lower than thegas temperature in the minihalo. Most of the gas will therefore be expelled from thehalo on a dynamical timescale. Dotted line: Adiabatic behavior. The gas in the outerreaches of the minihalo, as well as in the general IGM, approximately follows thisscaling. Grey symbols: Location of the gas that makes up the original stellar ejecta.
This gas is hot and diffuse, can consequently cool only slowly, and therefore provides
the pressure source that drives the expansion of the bubble. Notice that a fraction of thegas has been able to cool to ∼200 K due to the action of molecular hydrogen. (FromBromm et al. 2003.)

stars in the terminology of Mackey et al. (2003). This is consistent with theprediction in Mackey et al. (2003) that Population II.5 stars can only form in
DM halos massive enough to be able to cool via atomic hydrogen (see also
Salvaterra et al. 2003). To further test the possibility of SN-triggered star formation in the high-redshift universe, it will be necessary to explore a range ofcosmological environments corresponding to different halo masses and collapse
redshifts.
What is the fate of the metals that are produced in the PISN progenitor? PISNe
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0.5, and for the largest stellar masses in the PISN range, most of M Z is madeup of iron (Heger & Woosley 2002). Bromm et al. (2003) find that for E SN =1053 ergs, ∼90% of the metals have escaped from the DM minihalo 4 × 106
years after the explosion. The metals have filled most of the interior hot bubble(see inset of Figure 7), and there has not yet been sufficient time for them tosignificantly mix into the dense surrounding shell. A rough estimate of the resulting
metallicity in the surrounding contaminated IGM can be obtained as follows:
Z = M Z /Mgas ∼ 100 M /105 M  10−2 Z  . This value is well above the critical
metallicity threshold, Z crit ∼ 10−3.5 Z  , for enabling the formation of lower-massstars (see Section 3.5). The conclusion is that Population III stars can continue
to form in the pristine gas within the neighboring halos, as the metals that are
dispersed in the explosion do not reach them.
Whereas a single SN explosion might be able to evacuate most of the gas froma minihalo, the situation could be quite different for multiple SN explosions inearly dwarf galaxies of mass ∼108 M (Mori et al. 2002, Wada & Venkatesan
2003). In this case, the simulations indicate that off-nuclear SN explosions drive
inward-propagating blast waves that act to assemble gas in the central regions of
the host galaxy. The gas in these galaxies may thus be able to survive multiple SN
explosions, and star formation would not be significantly suppressed owing to themechanical SN feedback. In summary, it appears that the effect of SN events on
star formation in early galaxies and on the chemical and thermal evolution of theIGM depends sensitively on the specific environment.
For the PISN explosion studied by Bromm et al. (2003), the extent of the metalenriched region with a physical size of ∼1 kpc is comparable to the radius of therelic H II region around the minihalo. In a related paper (Yoshida et al. 2004),
it is argued that if Population III stars have led to an early partial reionizationof the universe, as may be required by the recent WMAP results (see Section5.1), this will have resulted in a nearly uniform enrichment of the universe to alevel Z min  10−4 Z  already at z  15. It would indeed be a remarkable featureof the universe if the first stars had endowed the IGM with such a widespread,near-universal level of early enrichment (see also Oh et al. 2001). An interesting
implication could be the observability of unsaturated metal absorption lines atz  6, which would probe the neutral fraction of the IGM during the epoch ofreionization (Oh 2002).
We end this section by emphasizing that we are only beginning to understandthe complex processes that distributed the first heavy elements into the IGM, andmore sophisticated numerical simulations are required for further progress.

5. OBSERVATIONAL SIGNATURE
An increasing variety of observational probes of the first stars has been suggested(e.g., Haiman & Loeb 1997, Oh et al. 2003, Stiavelli et al. 2004). Below, we focus
on three of them, and here only briefly mention some of the others (see Barkana

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107& Loeb 2001 for a comprehensive discussion). Intriguing possibilities would arise
if the Population III IMF were sufficiently top-heavy to enable the formation ofmassive black holes (e.g., Madau & Rees 2001, Islam et al. 2003, Volonteri et al.
2003). Possible consequences might include backgrounds of gravitational waves(Schneider et al. 2000, de Araujo et al. 2002) and diffuse neutrinos (Schneider
et al. 2002b). In the latter case, the high-energy (TeV) neutrinos might originate in
the relativistic jets that are associated with gamma-ray bursts (GRBs), which are
believed to accompany the birth of a black hole (see Section 5.3).
The integrated light emitted during the lifetime of Population III stars is predicted (Bond et al. 1986, Santos et al. 2002, Salvaterra & Ferrara 2003) to contributesignificantly to the locally observed cosmic near-infrared background (e.g., Hauser& Dwek 2001). A related signature is the Population III contribution to the background anisotropy (Magliocchetti et al. 2003, Cooray et al. 2004). Finally, 21-cm
radiation from neutral hydrogen, both in emission and absorption, could provide a
unique probe of the conditions in the early minihalos, thus testing our assumptions
on the initial conditions for Population III star formation (e.g., Madau et al. 1997;
Tozzi et al. 2000; Carilli et al. 2002; Furlanetto & Loeb 2002; Iliev et al. 2002,
2003; Ciardi & Madau 2003). However, in practice, these observations may be
extremely challenging owing to the strong foreground contamination (Di Matteo
et al. 2002).
In the following section, we discuss three empirical probes of the first stars insomewhat greater detail, as they have already provided us with constraints on thenature of the first stars, or promise soon to do so.

5.1. Reionization Signature from the First Stars
The WMAP satellite has recently measured, from the CMB polarization anisotropies (Kaplinghat et al. 2003), the total optical depth to Thomson scattering:τe = 0.17 ± 0.04 (Kogut et al. 2003, Spergel et al. 2003). Such a high value
is surprising. We know that the IGM was completely ionized again at redshiftsz  6 from the absence of Gunn-Peterson troughs in high-redshift Sloan Digital
Sky Survey (SDSS) quasars (e.g., Fan et al. 2003). Scattering from free electrons
between z ∼ 0–6 only contributes τe  0.05. The observed excess must have arisen
from the scattering of CMB photons off free electrons ionized by star formation
at higher redshifts. Normal stellar populations (characterized by a standard IMF,
or at most a modestly top-heavy one with characteristic mass Mc ∼ 5 M ) could
have successfully reproduced the measured high optical depth if reionization had
occurred already at z ∼ 15 (Ciardi et al. 2003).
Such a solution, however, runs into difficulties when two other empirical constraints are taken into account. The first one considers the thermal history of the
postreionization IGM (e.g., Theuns et al. 2002, Hui & Haiman 2003). After reionization is completed, the IGM continues to be heated by photoionization. This
heating, however, cannot compete with the cooling owing to adiabatic Hubble
expansion. In consequence, the earlier reionization occurred, the lower the IGM

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temperature at a given fixed redshift. Observations at z ∼ 3 indicate that thelast complete overlap of ionizing bubbles (i.e., the last time the IGM was completely ionized) could not have occurred before a redshift z ∼ 9. The second
constraint, which is difficult to accommodate within the model of Ciardi et al.
(2003), is the detection of select SDSS quasars at z  6 with Gunn-Peterson absorption in their vicinity. It is not easy to see how islands of neutral gas could
have persisted down to z ∼ 6 if the IGM had already been completely reionized
at z ∼ 15.
To simultaneously account for all the available observational constraints thus
poses a nontrivial challenge to theory. It has been suggested that an early generationof very massive Population III stars, as predicted by numerical simulations (seeSection 3), could naturally allow for a solution (Cen 2003a,b; Haiman & Holder
2003; Somerville & Livio 2003; Wyithe & Loeb 2003a,b). The high efficiency of
producing ionizing photons per stellar baryon for very massive Population III stars(e.g., Tumlinson & Shull 2000; Bromm et al. 2001b; Schaerer 2002, 2003) can leadto an early first episode of at least partial reionization at z ∼ 15, thus ensuring asufficiently large optical depth to Thomson scattering. As we discussed in Section
4, Population III star formation was self-limiting owing to the effect of radiationand the dispersal of heavy elements. After the era of Population III star formation
at z  15, the IGM was able to recombine again owing to the temporary lull in the
production of ionizing photons. Eventually, normal (Population II) stars form in
more massive (dwarf-sized) systems, and lead to the final, complete reionizationof the universe at z ∼ 7. In this case, there is no conflict with the empiricalconstraints outlined above. In Figure 8, a sample reionization history is shown thatis representative for the double-reionization models (from Wyithe & Loeb 2003b).
At present, the error bars in the WMAP determination of τe are still too large tofirmly establish the need for very massive Population III stars to have formed in
the high-redshift universe. By the end of its mission, however, the error on τe is
estimated to have gone down to ±0.01. If the high current value were to persist,the evidence for a top-heavy IMF in primordial star formation would have been
considerably strengthened.
Cosmological models with reduced small-scale power (see Section 2) cannot
easily be reconciled with a high value of τe (Somerville et al. 2003; Yoshida et al.
2003b,c). To compensate for the missing contribution from minihalos, one would
have to invoke an extremely high ionizing photon production rate in galaxies
at z  6. This could, in principle, be achieved by maintaining a top-heavy IMF
over a long period, 6  z  20. The problem with such a scenario is the severe
overproduction of metals. These would have to be hidden to an improbable degree
so that the formation of massive Population III stars were not terminated too early(Yoshida et al. 2004). The study of the first stars thus has the exciting potential toprovide astrophysical constraints on the properties of DM particles because theseproperties are reflected in the nature of the density fluctuations at small spatial
scales.

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Figure 8 Sample history of reionization. Shown are the volume filling factors (Q) of
ionized regions as a function of redshift. For hydrogen (solid line), a complex evolution
is evident, with an early (partial) reionization at z ∼ 15 owing to massive Population III
stars. This mode of star formation is self-terminating, such that the second (complete)reionization at z ∼ 9 is due to normal Population II stars. Also shown is the hypotheticalcase where He ionization is ignored (light lines), as well as the predicted histories for
He+ and He++ reionization. (Adapted from Wyithe & Loeb 2003b.)5.2. Stellar Archaeology: Relics from the End of the Dark Ages
It has long been realized that the most metal-poor stars found in our cosmic neighborhood would encode the signature from the first stars within their elemental
abundance pattern (e.g., Bond 1981, Beers et al. 1992). For many decades, however, the observational search has failed to discover a truly first-generation starwith zero metallicity. Indeed, there seemed to have been an observational lower
limit of [Fe/H] ∼ − 4 (e.g., Carr 1987, Oey 2003). In view of the recent theoreticalprediction that most Population III stars were very massive, with associated lifetimes of ∼106 years, the failure to find any living Population III star in the Galaxy
is not surprising, as they all would have died a long time ago (e.g., Hernandez& Ferrara 2001). Furthermore, theory has predicted that star formation out of extremely low-metallicity gas, with Z  Z crit ∼ 10−3.5 Z  (see Section 3.5), would
be essentially equivalent to that out of truly primordial gas. Again, this theoreticalprediction was in accordance with the apparent observed lower-metallicity cutoff.
Recently, however, this simple picture has been challenged by the discovery of
the star HE0107-5240 with a mass of 0.8 M and an iron abundance of [Fe/H] =−5.3 (Christlieb et al. 2002). This finding indicates that at least some low mass
stars could have formed out of extremely low-metallicity gas. Does the existence of
this star invalidate the theory of a metallicity threshold for enabling low-mass starformation? As pointed out by Umeda & Nomoto (2003), a possible explanation

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Figure 9 Observed abundances in low-metallicity galactic halo stars. For carbon(upper panel) and oxygen (lower panel), filled circles correspond to samples of dwarf
and subgiant stars (from Akerman et al. 2003), and open squares correspond to asample of giant stars (from Cayrel et al. 2004). The dashed lines indicate the predicted
critical carbon and oxygen abundances (see Section 3.5). The cross is the location of theextremely iron-poor giant star HE0107-5240. (Adapted from Bromm & Loeb 2003c.)

could lie in the unusually high abundances of carbon and oxygen in HE01075240. As discussed in Section 3.5, Bromm & Loeb (2003c) have demonstrated
that indeed carbon and oxygen may have been the main drivers for the Population
III to II transition.
In Figure 9, we compare the theoretical thresholds derived by Bromm & Loeb
(2003c) to the observed carbon and oxygen abundances in metal-poor dwarf(Akerman et al. 2003) and giant (Cayrel et al. 2004) stars in the halo of our
Galaxy. All data points lie above the critical oxygen abundance, but a few cases lie
below the critical carbon threshold. All these low-mass stars are consistent withthe model because the corresponding oxygen abundances lie above the predicted
threshold. The subcritical [C/H] abundances could have originated either in the5 Aug 2004 22:8

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progenitor cloud or from the mixing of CNO-processed material (with carbon
converted into nitrogen) into the stellar atmosphere during the red giant phase.
To guard against this a posteriori processing of carbon (and to a lesser degree ofoxygen), dwarf stars should be used, as they provide a more reliable record of theprimordial abundance pattern that was present at birth. The current data on dwarf
stars (filled symbols in Figure 9) does not yet reach carbon and oxygen abundancesthat are sufficiently low to probe the theoretical cooling thresholds. Note also thatthe extremely iron-poor star HE0107-5240 has carbon and oxygen abundancesthat lie above the respective critical levels. The formation of this low-mass star(∼0.8 M ) is therefore consistent with the theoretical framework considered byBromm & Loeb (2003c).
The lessons from stellar archaeology on the nature of the first stars are likely
to increase in importance because greatly improved, large surveys of metal-poorgalactic halo stars are under way or are currently being planned.

5.3. Gamma-Ray Bursts as Probes of the First Stars
GRBs are the brightest electromagnetic explosions in the universe, and they should
be detectable out to very high redshifts (e.g., Wijers et al. 1998, Blain & Natarajan
2000, Ciardi & Loeb 2000, Lamb & Reichart 2000, Bromm & Loeb 2002,
Choudhury & Srianand 2002). Although the nature of the central engine that powers the relativistic jets is still debated, recent evidence indicates that GRBs trace
the formation of massive stars (e.g., Bloom et al. 2002, and references therein).
Because the first stars are predicted to be predominantly very massive, their deathmight possibly give rise to GRBs at very high redshifts (e.g., Schneider et al.
2002b). A detection of the highest-redshift GRBs would probe the earliest epochsof star formation, one massive star at a time (e.g., Barkana & Loeb 2004, Mészáros& Rees 2003). The Swift satellite,4 planned for launch in 2004, is expected to detectapproximately one hundred GRBs per year. The redshifts of high-z GRBs can beeasily measured through infrared photometry on the basis of the Gunn-Peterson
trough in their spectra owing to Lyα absorption by neutral intergalactic hydrogen along the line of sight. What fraction of the detected bursts will originate atredshifts z  5?
To assess the utility of GRBs as probes of the first stars, one has to calculate theexpected redshift distribution of GRBs (e.g., Bromm & Loeb 2002, Choudhury &Srianand 2002). Under the assumption that the GRB rate is simply proportional to
the star formation rate, Bromm & Loeb (2002) find that approximately a quarterof all GRBs detected by Swift will originate from a redshift z  5 (see Figure 10).
This estimate takes into account the detector sensitivity, but it is still very uncertain
because of the poorly determined GRB luminosity function. The rate of highredshift GRBs may, however, be significantly suppressed if the early massive stars
4

See http://swift.gsfc.nasa.gov.

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Figure 10 Redshift distribution of all GRBs in comparison with measurements influx-limited surveys. (a) Fraction of bursts that originate at a redshift higher than z
versus z. The data points reflect ∼20 observed redshifts to date. (b) Fraction of bursts
per logarithmic interval of (1 + z) versus z. (Solid lines) All GRBs for star formation inhalos massive enough to allow cooling via lines of atomic hydrogen. The calculation
assumes that the GRB rate is proportional (with a constant factor) to the star formation
rate at all redshifts. (Dot-dashed lines) Expected distribution for Swift. (Long-dashedlines) Expected distribution for BATSE. The curves for the two flux-limited surveys
are rather uncertain because of the poorly determined GRB luminosity function. (FromBromm & Loeb 2002.)

fail to launch a relativistic outflow. This is conceivable, as metal-free stars may
experience negligible mass loss before exploding as a SN (Baraffe et al. 2001).
They would then retain their massive hydrogen envelope, and any relativistic jet
might be quenched before escaping from the star (Heger et al. 2003).
If high-redshift GRBs exist, the launch of Swift will open up an exciting new
window into the cosmic dark ages (e.g., Gou et al. 2004). Different from quasars
or galaxies that fade with increasing redshift, GRB afterglows maintain a roughly
constant observed flux at different redshifts for a fixed observed time lag after the
gamma-ray trigger (Ciardi & Loeb 2000). The increase in the luminosity distanceat higher redshifts is compensated by the fact that a fixed observed time lag corresponds to an intrinsic time shorter by a factor of (1 + z) in the source rest-frame,during which the GRB afterglow emission is brighter. This quality makes GRB
afterglows the best probes of the metallicity and ionization state of the intervening
IGM during the epoch of reionization. In contrast to quasars, the UV emission from

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GRBs has a negligible effect on the surrounding IGM (because ∼1051 ergs can only
ionize ∼4 × 104 M of hydrogen). Moreover, the host galaxies of GRBs induce amuch weaker perturbation to the Hubble flow in the surrounding IGM, comparedto the massive hosts of the brightest quasars (Barkana & Loeb 2004). Hence, GRB
afterglows offer the ideal probe (much better than quasars or bright galaxies) of
the damping wing of the Gunn-Peterson trough (Miralda-Escudé 1998) that signals the neutral fraction of the IGM as a function of redshift during the epoch ofreionization.

6. EPILOGUE: THE ROAD AHEAD
One of the fascinating aspects of the recent research on the first stars has been theconfluence of many hitherto unrelated subfields. To fully explore its rich potential,
we have to utilize high-redshift studies and near-field cosmological data (Freeman& Bland-Hawthorn 2002). The latter approach refers to the stellar archaeology ofobserving the chemical abundance patterns in extremely metal-poor galactic halostars (see Section 5.2). Furthermore, the study of the first stars lies at the intersectionof star-formation theory with cosmological structure formation, where the scales
for the phenomena overlap. Similarly, there is also an overlap between star andgalaxy formation because at this early time, the scale of galaxies and the scaleof star-forming clouds had not yet become clearly separated. Therefore, we caninvestigate all these phenomena in their most elementary form, before their scales
had become widely separated and before the universe had evolved into the vastly
more complicated structure that we observe today, where there is a huge range ofscales at play.
Most of the work so far has been theoretical, but the field is driven by the prospectof upcoming observational probes that will soon test the theoretical predictions.
The WMAP satellite has already provided important clues, and Swift will soon doso as well. Most importantly, NASA’s planned successor mission to the Hubble
Space Telescope, the JWST, to be launched around 2011, is a prime driver behindcurrent cosmological research. The JWST is designed to image the first stars,
SNe, and quasars. In thinking ahead, pondering about the first glimpses that wehave gained into the cosmic dark ages, one wonders what exciting discoveries areawaiting us.
ACKNOWLEDGMENTS
We express our gratitude to Paolo Coppi for many discussions on the work discussed in this review. V.B. thanks Avi Loeb and Lars Hernquist at the HarvardSmithsonian Center for Astrophysics for support from NSF grant AST 00-71019,
as well as Cathie Clarke at the Institute of Astronomy in Cambridge from theEuropean Community’s Research Training Network under contract HPRN-CT2000-0155, “Young Stellar Clusters.”

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http://astro.annualreviews.org

LITERATURE CITED
Abel T, Anninos P, Norman ML, Zhang Y. 1998.
Ap. J. 508:518–29
Abel T, Anninos P, Zhang Y, Norman ML. 1997.
New Astron. 2:181–207
Abel T, Bryan G, Norman ML. 2000. Ap. J.
540:39–44
Abel T, Bryan G, Norman ML. 2002. Science
295:93–98
Aguirre A, Hernquist L, Schaye J, Katz N,
Weinberg DH, Gardner J. 2001a. Ap. J.
561:521–49
Aguirre A, Hernquist L, Schaye J, Weinberg
DH, Katz N, Gardner J. 2001b. Ap. J.
560:599–605
Akerman CJ, Carigi L, Nissen PE, Pettini
M, Asplund M. 2003. Astron. Astrophys.
414:931–42
Anninos P, Norman ML. 1996. Ap. J. 460:556–
68
Baade W. 1944. Ap. J. 100:137–46
Balsara D. 2001. J. Korean Astron. Soc.
34:S181–90
Baraffe I, Heger A, Woosley SE. 2001. Ap. J.
550:890–96
Barkana R, Loeb A. 1999. Ap. J. 523:54–
65
Barkana R, Loeb A. 2001. Phys. Rep. 349:125–
238
Barkana R, Loeb A. 2004. Ap. J. 601:64–
77
Barkat Z, Rakavy G, Sack N. 1967. Phys. Rev.
Lett. 18:379–81
Bate MR, Bonnell IA, Bromm V. 2002. MNRAS
332:L65–68
Bate MR, Bonnell IA, Bromm V. 2003. MNRAS
339:577–99
Bate MR, Bonnell IA, Price NM. 1995. MNRAS
277:362–76
Beers TC, Preston GW, Shectman SA. 1992.
Astron. J. 103:1987–2034
Blain AW, Natarajan P. 2000. MNRAS 312:
L35–38Bloom JS, Kulkarni SR, Djorgovski SG. 2002.
Astron. J. 123:1111–48
Bode P, Ostriker JP, Turok N. 2001. Ap. J.
556:93–107
Bond HE. 1981. Ap. J. 248:606–11
Bond JR, Arnett WD, Carr BJ. 1984. Ap. J.
280:825–47
Bond JR, Carr BJ, Hogan CJ. 1986. Ap. J.
306:428–50
Bromm V. 2000. Star formation in the early universe. PhD thesis. Yale Univ., New Haven,CT. 165 pp.
Bromm V, Clarke CJ. 2002. Ap. J. 566:L1–4
Bromm V, Coppi PS, Larson RB. 1999. Ap. J.
527:L5–8
Bromm V, Coppi PS, Larson RB. 2002. Ap. J.
564:23–51
Bromm V, Ferrara A, Coppi PS, Larson RB.
2001a. MNRAS 328:969–76
Bromm V, Kudritzki RP, Loeb A. 2001b. Ap. J.
552:464–72
Bromm V, Loeb A. 2002. Ap. J. 575:111–16
Bromm V, Loeb A. 2003a. Ap. J. 596:34–46
Bromm V, Loeb A. 2003b. In The Emergence of
Cosmic Structure, ed. SS Holt, CS Reynolds,
pp. 73–83. Woodbury, NY: Am. Inst. Phys.
Bromm V, Loeb A. 2003c. Nature 425:812–14
Bromm V, Loeb A. 2004. New Astron. 9:353–64
Bromm V, Yoshida N, Hernquist L. 2003. Ap.
J. 596:L135–38
Carilli CL, Gnedin NY, Owen F. 2002. Ap. J.
577:22–30
Carlberg RG. 1981. MNRAS 197:1021–29
Carr BJ. 1987. Nature 326:829–30
Carr BJ. 1994. Annu. Rev. Astron. Astrophys.
32:531–90
Carr BJ, Bond JR, Arnett WD. 1984. Ap. J.
277:445–69
Cayrel R. 1986. Astron. Astrophys. 168:81–88
Cayrel R, Depagne E, Spite M, Hill V, Spite F,
et al. 2004. Astron. Astrophys. 416:1117–38
Cen R. 2003a. Ap. J. 591:L5–8

5 Aug 2004 22:8

AR

AR222-AA42-03.tex

AR222-AA42-03.sgm

LaTeX2e(2002/01/18)

FIRST STARS
Cen R. 2003b. Ap. J. 591:12–37
Choudhury TR, Srianand R. 2002. MNRAS
336:L27–31
Christlieb N, Bessell MS, Beers TC, Gustafsson
B, Korn A, et al. 2002. Nature 419:904–6
Ciardi B, Ferrara A, Abel T. 2000a. Ap. J.
533:594–600
Ciardi B, Ferrara A, Governato F, Jenkins A.
2000b. MNRAS 314:611–29
Ciardi B, Ferrara A, White SDM. 2003. MNRAS
344:L7–11
Ciardi B, Loeb A. 2000. Ap. J. 540:687–96
Ciardi B, Madau P. 2003. Ap. J. 596:1–8
Clarke CJ, Bromm V. 2003. MNRAS 343:1224–
30
Cooray A, Bock JJ, Keating B, Lange AE,
Matsumoto T. 2004. Ap. J. In press (astroph/0308407)
Couchman HMP, Rees MJ. 1986. MNRAS
221:53–62
de Araujo JCN, Miranda O, Aguiar OD. 2002.
MNRAS 330:651–59
Dekel A, Silk J. 1986. Ap. J. 303:39–55
Dijkstra M, Haiman Z, Rees MJ, Weinberg DH.
2004. Ap. J. 601:666–75
Di Matteo T, Perna R, Abel T, Rees MJ. 2002.
Ap. J. 564:576–80
Draine BT, Bertoldi F. 1996. Ap. J. 468:269–89
Eisenstein DJ, Loeb A. 1995. Ap. J. 443:11–
17
Elmegreen BG. 1994. Ap. J. 427:384–87
Fan X, Strauss MA, Schneider DP, Becker RH,
White RL, et al. 2003. Astron. J. 125:1649–
59
Ferrara A. 1998. Ap. J. 499:L17–20
Ferrara A, Pettini M, Shchekinov Y. 2000.
MNRAS 319:539–48
Flores RA, Primack JR. 1994. Ap. J. 427:L1–4
Flower DR, Le Bourlot J, Pineau des Forêts G,
Roueff E. 2000. MNRAS 314:753–58
Freeman K, Bland-Hawthorn J. 2002. Annu.
Rev. Astron. Astrophys. 40:487–537
Fryer CL, Woosley SE, Heger A. 2001. Ap. J.
550:372–82
Fuller TM, Couchman HMP. 2000. Ap. J.
544:6–20
Furlanetto SR, Loeb A. 2002. Ap. J. 579:1–9
Furlanetto SR, Loeb A. 2003. Ap. J. 588:18–34

P1: IKH

115

Galli D, Palla F. 1998. Astron. Astrophys.
335:403–20
Glover SCO, Brand PWJL. 2001. MNRAS
321:385–97
Glover SCO, Brand PWJL. 2003. MNRAS
340:210–26
Gou LJ, Mészáros P, Abel T, Zhang B. 2004.
Ap. J. 604:508–20
Gnedin NY, Ostriker JP. 1997. Ap. J. 486:581–
98
Haiman Z, Abel T, Rees MJ. 2000. Ap. J.
534:11–24
Haiman Z, Holder GP. 2003. Ap. J. 595:1–12
Haiman Z, Loeb A. 1997. Ap. J. 483:21–37
Haiman Z, Loeb A. 2001. Ap. J. 552:459–63
Haiman Z, Rees MJ, Loeb A. 1997. Ap. J.
476:458–63. Erratum. 1997. Ap. J. 484:985
Haiman Z, Thoul AA, Loeb A. 1996. Ap. J.
464:523–38
Hauser MG, Dwek E. 2001. Annu. Rev. Astron.
Astrophys. 39:249–307
Heger A, Fryer CL, Woosley SE, Langer N,
Hartmann DH. 2003. Ap. J. 591:288–300
Heger A, Woosley SE. 2002. Ap. J. 567:532–43
Hernandez X, Ferrara A. 2001. MNRAS
324:484–90
Hirashita H, Ferrara A. 2002. MNRAS 337:921–
37
Holder GP, Haiman Z, Kaplinghat M, Knox L.
2003. Ap. J. 595:13–18
Hollenbach D, McKee CF. 1979. Ap. J. Suppl.
41:555–92
Hollenbach D, McKee CF. 1989. Ap. J.
342:306–36
Hui L, Haiman Z. 2003. Ap. J. 596:9–18
Hutchings RM, Santoro F, Thomas PA, Couchman HMP. 2002. MNRAS 330:927–36
Hutchins JB. 1976. Ap. J. 205:103–21
Iliev IT, Scannapieco E, Martel H, Shapiro PR.
2003. MNRAS 341:81–90
Iliev IT, Shapiro PR, Ferrara A, Martel H. 2002.
Ap. J. 572:L123–26
Islam RR, Taylor JE, Silk J. 2003. MNRAS
340:647–56
Jang-Condell H, Hernquist L. 2001. Ap. J.
548:68–78
Jeans JH. 1902. Phil. Trans. Roy. Soc. 199:1–53
Kaiser N. 1984. Ap. J. 284:L9–12

5 Aug 2004 22:8

116

AR

AR222-AA42-03.tex

AR222-AA42-03.sgm

LaTeX2e(2002/01/18)

P1: IKH

BROMM  LARSON

Kaplinghat M, Chu M, Haiman Z, Holder G,
Knox L, Skordis C. 2003. Ap. J. 583:24–32
Kashlinsky A, Rees MJ. 1983. MNRAS
205:955–71
Katz N. 1991. Ap. J. 368:325–36
Kirshner RP. 2003. Science 300:1914–18
Kitayama T, Susa H, Umemura M, Ikeuchi S.
2001. MNRAS 326:1353–66
Kitsionas S, Whitworth AP. 2002. MNRAS
330:129–36
Kogut A, Spergel DN, Barnes C, Bennett
CL, Halpern M, et al. 2003. Ap. J. Suppl.
148:161–73
Kroupa P. 2002. Science 295:82–91
Kulsrud RM, Cen R, Ostriker JP, Ryu D. 1997.
Ap. J. 480:481–91
Lamb DQ, Reichart DE. 2000. Ap. J. 536:1–18
Larson RB. 1969. MNRAS 145:271–95
Larson RB. 1974. MNRAS 169:229–45
Larson RB. 1998. MNRAS 301:569–81
Larson RB. 2002. MNRAS 332:155–64
Larson RB. 2003. Rep. Prog. Phys. 66:1651–97
Lepp S, Shull JM. 1984. Ap. J. 280:465–69
Loeb A, Barkana R. 2001. Annu. Rev. Astron.
Astrophys. 39:19–66
Loeb A, Haiman Z. 1997. Ap. J. 490:571–76
Loeb A, Rasio FA. 1994. Ap. J. 432:52–61
Low C, Lynden-Bell D. 1976. MNRAS
176:367–90
Machacek ME, Bryan G, Abel T. 2001. Ap. J.
548:509–21
Machacek ME, Bryan G, Abel T. 2003. MNRAS
338:273–86
Mackey J, Bromm V, Hernquist L. 2003. Ap. J.
586:1–11
Mac Low MM, Shull JM. 1986. Ap. J. 302:585–
89
Madau P, Ferrara A, Rees MJ. 2001. Ap. J.
555:92–105
Madau P, Meiksin A, Rees MJ. 1997. Ap. J.
475:429–44
Madau P, Rees MJ. 2001. Ap. J. 551:L27–30
Madau P, Rees MJ, Volonteri M, Haardt F, Oh
SP. 2004. Ap. J. 604:484–94
Magliocchetti M, Salvaterra R, Ferrara A. 2003.
MNRAS 342:L25–29
Matsuda T, Sato H, Takeda H. 1969. Prog.
Theor. Phys. 42:219–33

McDowell MRC. 1961. Observatory 81:240–43
Mészáros P, Rees MJ. 2003. Ap. J. 591:L91–
94
Miralda-Escudé J. 1998. Ap. J. 501:15–22
Miralda-Escudé J. 2003. Science 300:1904–9
Monaghan JJ. 1992. Annu. Rev. Astron. Astrophys. 30:543–74
Moore B, Ghigna S, Governato F, Lake G,
Quinn T, et al. 1999. Ap. J. 524:L19–22
Mori M, Ferrara A, Madau P. 2002. Ap. J.
571:40–55
Motte F, André P, Neri R. 1998. Astron. Astrophys. 336:150–72
Nakamura F, Umemura M. 1999. Ap. J.
515:239–48
Nakamura F, Umemura M. 2001. Ap. J. 548:19–
32
Nakamura F, Umemura M. 2002. Ap. J.
569:549–57
Nishi R, Tashiro M. 2000. Ap. J. 537:50–54
Nozawa T, Kozasa T, Umeda H, Maeda K,
Nomoto K. 2003. Ap. J. 598:785–803
Ober WW, El Eid MF, Fricke KJ. 1983. Astron.
Astrophys. 119:61–68
Oey MS. 2003. MNRAS 339:849–60
Oh SP. 2001. Ap. J. 553:499–512
Oh SP. 2002. MNRAS 336:1021–29
Oh SP, Cooray A, Kamionkowski M. 2003.
MNRAS 342:L20–24
Oh SP, Haiman Z. 2002. Ap. J. 569:558–72
Oh SP, Haiman Z. 2003. MNRAS 346:456–
72
Oh SP, Nollett KM, Madau P, Wasserburg GJ.
2001. Ap. J. 562:L1–4
Omukai K. 2000. Ap. J. 534:809–24
Omukai K. 2001. Ap. J. 546:635–51
Omukai K, Inutsuka S. 2002. MNRAS 332:59–
64
Omukai K, Nishi R. 1998. Ap. J. 508:141–50
Omukai K, Nishi R. 1999. Ap. J. 518:64–68
Omukai K, Palla F. 2001. Ap. J. 561:L55–58
Omukai K, Palla F. 2003. Ap. J. 589:677–87
Omukai K, Yoshii Y. 2003. Ap. J. 599:746–58
Ostriker JP, Gnedin NY. 1996. Ap. J. 472:L63–
67
Ostriker JP, McKee CF. 1988. Rev. Mod. Phys.
60:1–68

5 Aug 2004 22:8

AR

AR222-AA42-03.tex

AR222-AA42-03.sgm

LaTeX2e(2002/01/18)

FIRST STARS
Ostriker JP, Steinhardt P. 2003. Science
300:1909–13
Padoan P, Nordlund A. 2002. Ap. J. 576:870–
79
Palla F, Salpeter EE, Stahler SW. 1983. Ap. J.
271:632–41
Peebles PJE, Dicke RH. 1968. Ap. J. 154:891–
908
Peiris HV, Komatsu E, Verde L, Spergel DN,
Bennett CL, et al. 2003. Ap. J. Suppl.
148:213–31
Penston MV. 1969. MNRAS 144:425–48
Pudritz RE. 2002. Science 295:68–76
Puy D, Signore M. 1996. Astron. Astrophys.
305:371–78
Puy D, Signore M. 1999. New Astron. Rev.
43:223–41
Qian Y-Z, Sargent WLW, Wasserburg GJ. 2002.
Ap. J. 569:L61–64
Qian Y-Z, Wasserburg GJ. 2002. Ap. J.
567:515–31
Rees MJ. 1976. MNRAS 176:483–86
Rees MJ. 1993. Quart. J. Roy. Astron. Soc.
34:279–89
Rees MJ, Ostriker JP. 1977. MNRAS 179:541–
59
Ricotti M, Gnedin NY, Shull JM. 2001. Ap. J.
560:580–91
Ricotti M, Gnedin NY, Shull JM. 2002a. Ap. J.
575:33–48
Ricotti M, Gnedin NY, Shull JM. 2002b. Ap. J.
575:49–67
Ripamonti E, Haardt F, Ferrara A, Colpi M.
2002. MNRAS 334:401–18
Salpeter EE. 1955. Ap. J. 121:161–67
Salvaterra R, Ferrara A. 2003. MNRAS
339:973–82
Santoro F, Thomas PA. 2003. MNRAS
340:1240–48
Santos MR, Bromm V, Kamionkowski M.
2002. MNRAS 336:1082–92
Sargent WLW. 1986. In The Lesson of Quantum
Theory, ed. J de Boer, E Dal, O Ulfbeck, pp.
263–82. Amsterdam: Elsevier
Saslaw WC, Zipoy D. 1967. Nature 216:976–
78
Scannapieco E, Ferrara A, Madau P. 2002. Ap.
J. 574:590–98

P1: IKH

117

Scannapieco E, Schneider R, Ferrara A. 2003.
Ap. J. 589:35–52
Schaerer D. 2002. Astron. Astrophys. 382:28–
42
Schaerer D. 2003. Astron. Astrophys. 397:527–
38
Schneider R, Ferrara A, Ciardi B, Ferrari V,
Matarrese S. 2000. MNRAS 317:385–90
Schaye J, Aguirre A, Kim T-S, Theuns T, Rauch
M, Sargent WLW. 2003. Ap. J. 596:768–
96
Schneider R, Ferrara A, Natarajan P, Omukai
K. 2002a. Ap. J. 571:30–39
Schneider R, Ferrara A, Salvaterra R, Omukai
K, Bromm V. 2003a. Nature 422:869–71
Schneider R, Ferrara A, Salvaterra R. 2003b.
MNRAS In press (astro-ph/0307087)
Schneider R, Guetta D, Ferrara A. 2002b.
MNRAS 334:173–81
Schwarzschild M, Spitzer L. 1953. Observatory
73:77–79
Seager S, Sasselov D, Scott D. 2000. Ap. J.
Suppl. 128:407–30
Shapiro PR, Iliev IT, Raga AC. 2004. MNRAS
348:753–82
Shapiro PR, Kang H. 1987. Ap. J. 318:32–65
Silk J. 1977. Ap. J. 211:638–48
Silk J. 1983. MNRAS 205:705–18
Sneden CS, Cowan JJ. 2003. Science 299:70–
75
Sokasian A, Abel T, Hernquist L, Springel V.
2003a. MNRAS 344:607–24
Sokasian A, Yoshida N, Abel T, Hernquist L,
Springel V. 2003b. MNRAS In press (astroph/0307451)
Somerville RS, Bullock JS, Livio M. 2003. Ap.
J. 593:616–21
Somerville RS, Livio M. 2003. Ap. J. 593:611–
15
Songaila A. 2001. Ap. J. 561:L153–6. Erratum.
2002. Ap. J. 568:L139
Spergel DN, Verde L, Peiris HV, Komatsu E,
Nolta MR, et al. 2003. Ap. J. Suppl. 148:175–
94
Stahler SW, Palla F, Salpeter EE. 1986. Ap. J.
302:590–605
Stecher TP, Williams DA. 1967. Ap. J.
149:L29–30

5 Aug 2004 22:8

118

AR

AR222-AA42-03.tex

AR222-AA42-03.sgm

LaTeX2e(2002/01/18)

P1: IKH

BROMM  LARSON

Stiavelli M, Fall SM, Panagia N. 2004. Ap. J.
600:508–19
Susa H, Umemura M. 2000. Ap. J. 537:578–
88
Tan JC, Blackman EG. 2004. Ap. J. 603:401–
13
Tan JC, McKee CF. 2004. Ap. J. 603:383–400
Tassis K, Abel T, Bryan G, Norman ML. 2003.
Ap. J. 587:13–24
Tegmark M, Silk J, Rees MJ, Blanchard A, Abel
T, Palla F. 1997. Ap. J. 474:1–12
Thacker RJ, Scannapieco E, Davis M. 2002. Ap.
J. 581:836–43
Theuns T, Schaye J, Zaroubi S, Kim T-S,
Tzanavaris P, Carswell B. 2002. Ap. J.
567:L103–6
Todini P, Ferrara A. 2001. MNRAS 325:726–36
Tozzi P, Madau P, Meiksin A, Rees MJ. 2000.
Ap. J. 528:597–606
Tumlinson J, Shull JM. 2000. Ap. J. 528:L65–
68
Tumlinson J, Shull JM, Venkatesan A. 2003.
Ap. J. 584:608–20
Uehara H, Inutsuka S. 2000. Ap. J. 531:L91–94
Uehara H, Susa H, Nishi R, Yamada M, Nakamura T. 1996. Ap. J. 473:L95–98
Umeda H, Nomoto K. 2002. Ap. J. 565:385–
404
Umeda H, Nomoto K. 2003. Nature 422:871–
73
Umemura M, Loeb A, Turner EL. 1993. Ap. J.
419:459–68
Venkatesan A, Giroux ML, Shull JM. 2001. Ap.
J. 563:1–8

Venkatesan A, Truran JW. 2003. Ap. J. 594:L1–
4Venkatesan A, Tumlinson J, Shull JM. 2003.
Ap. J. 584:621–32
Volonteri M, Haardt F, Madau P. 2003. Ap. J.
582:559–73
Wada K, Venkatesan A. 2003. Ap. J. 591:38–
42
Ward-Thompson D. 2002. Science 295:76–81
Whalen D, Abel T, Norman ML. 2003. Ap. J.
In press (astro-ph/0310283)
Wijers RAMJ, Bloom JS, Bagla JS, Natarajan
P. 1998. MNRAS 294:L13–17
Wolfire MG, Cassinelli JP. 1987. Ap. J.
319:850–67
Wood K, Loeb A. 2000. Ap. J. 545:86–99
Wyithe JSB, Loeb A. 2003a. Ap. J. 586:693–
708
Wyithe JSB, Loeb A. 2003b. Ap. J. 588:L69–72
Yoneyama T. 1972. Publ. Astron. Soc. Jpn.
24:87–98
Yoshida N, Abel T, Hernquist L, Sugiyama N.
2003a. Ap. J. 592:645–63
Yoshida N, Bromm V, Hernquist L. 2004. Ap.
J. 605:579–90
Yoshida N, Sokasian A, Hernquist L, Springel
V. 2003b. Ap. J. 591:L1–4
Yoshida N, Sokasian A, Hernquist L, Springel
V. 2003c. Ap. J. 598:73–85
Yoshida N, Sugiyama N, Hernquist L. 2003d.
MNRAS 344:481–91
Yoshii Y, Sabano Y. 1979. Publ. Astron. Soc.
Jpn. 31:505–21
Yoshii Y, Saio H. 1986. Ap. J. 301:587–600

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Figure 1 Emergence of primordial star-forming regions within a standard CDM
cosmology. Shown is the projected gas density at z  17 within a simulation box ofphysical size ~50 kpc. The bright knots at the intersections of the filamentary network are the sites where the first stars form. (From Yoshida et al. 2003a.)

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BROMM ■ LARSON

Figure 7 Situation ~106 years after the explosion for the case with ESN  1053 ergs.
Shown is the projected gas density within a box of linear physical size 1 kpc. The
SN bubble has expanded to a radius of ~200 pc, having evacuated most of the gas in
the minihalo (with a virial radius of ~150 pc). The dense shell has fragmented into
numerous cloudlets. Inset: Metal distribution after 3 Myr. The stellar ejecta (red dots)
trace the metals, and are embedded in pristine, un-enriched gas (blue dots). Morethan 90% of the heavy elements have escaped the minihalo and have filled a significant fraction of the hot bubble. (From Bromm et al. 2003.)