ARlogo Annu. Rev. Astron. Astrophys. 2012. 50:531-608
Copyright © 2012 by Annual Reviews. All rights reserved

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1.1. Introduction

Star formation encompasses the origins of stars and planetary systems, but it is also a principal agent of galaxy formation and evolution, and hence a subject at the roots of astrophysics on its largest scales.

The past decade has witnessed an unprecedented stream of new observational information on star formation on all scales, thanks in no small part to new facilities such as the Galaxy Evolution Explorer (GALEX), the Spitzer Space Telescope, the Herschel Space Observatory, the introduction of powerful new instruments on the Hubble Space Telescope (HST), and a host of groundbased optical, infrared, submillimeter, and radio telescopes. These new observations are providing a detailed reconstruction of the key evolutionary phases and physical processes that lead to the formation of individual stars in interstellar clouds, while at the same time extending the reach of integrated measurements of star formation rates (SFRs) to the most distant galaxies known. The new data have also stimulated a parallel renaissance in theoretical investigation and numerical modelling of the star formation process, on scales ranging from individual protostellar and protoplanetary systems to the scales of molecular clouds and star clusters, entire galaxies and ensembles of galaxies, even to the first objects, which are thought to have reionized the Universe and seeded today's stellar populations and Hubble sequence of galaxies.

This immense expansion of the subject, both in terms of the volume of results and the range of physical scales explored, may help to explain one of its idiosyncracies, namely the relative isolation between the community studying individual star-forming regions and stars in the Milky Way (usually abbreviated hereafter as MW, but sometimes referred to as "the Galaxy"), and the largely extragalactic community that attempts to characterise the star formation process on galactic and cosmological scales. Some aspects of this separation have been understandable. The key physical processes that determine how molecular clouds contract and fragment into clumps and cores and finally clusters and individual stars can be probed up close only in the Galaxy, and much of the progress in this subject has come from in-depth case studies of individual star-forming regions. Such detailed observations have been impossible to obtain for even relatively nearby galaxies. Instead the extragalactic branch of the subject has focused on the collective effects of star formation, integrated over entire star-forming regions, or often over entire galaxies. It is this collective conversion of baryons from interstellar gas to stars and the emergent radiation and mechanical energy from the stellar populations that is most relevant to the formation and evolution of galaxies. As a result, much of our empirical knowledge of star formation on these scales consists of scaling laws and other parametric descriptions, in place of a rigorous, physically-based characterization. Improving our knowledge of large-scale star formation and its attendant feedback processes is essential to understanding the birth and evolution of galaxies.

Over the past several years, it has become increasingly clear that many of the key processes influencing star formation on all of these scales lie at the interface between the scales within individual molecular clouds and those of galactic disks. It is now clear that the large scale SFR is determined by a hierarchy of physical processes spanning a vast range of physical scales: the accretion of gas onto disks from satellite objects and the intergalactic medium (Mpc); the cooling of this gas to form a cool neutral phase (kpc); the formation of molecular clouds (~ 10 - 100 pc); the fragmentation and accretion of this molecular gas to form progressively denser structures such as clumps (~ 1 pc) and cores (~ 0.1 pc); and the subsequent contraction of the cores to form stars (R) and planets (~ AU). The first and last of these processes operate on galactic (or extragalactic) and local cloud scales, respectively, but the others occur at the boundaries between these scales, and the coupling between processes is not yet well understood. Indeed it is possible that different physical processes provide the "critical path" to star formation in different interstellar and galactic environments. Whatever the answer to these challenging questions, however, Nature is strongly signalling that we need a unified approach to understanding star formation, one which incorporates observational and astrophysical constraints on star formation efficiencies, mass functions, etc. from small-scale studies along with a much deeper understanding of the processes that trigger and regulate the formation of star-forming clouds on galactic scales, which, in turn, set the boundary conditions for star formation within clouds.

This review makes a modest attempt to present a consolidated view of large-scale star formation, one which incorporates our new understanding of the star formation within clouds and the ensemble properties of star formation in our own MW into our more limited but broader understanding of star formation in external galaxies. As will become clear, the subject itself is growing and transforming rapidly, so our goal is to present a progress report in what remains an exciting but relatively immature field. Nevertheless a number of factors make this a timely occasion for such a review. The advent of powerful multi-wavelength observations has transformed this subject in fundamental ways since the last large observational review of the galaxy-scale aspects (Kennicutt 1998a), hereafter denoted K98. Reviews have covered various observational aspects of star formation within the MW (Evans 1999; Bergin & Tafalla 2007; Zinnecker & Yorke 2007), but the most comprehensive reviews have been theoretically based (Shu et al. 1987; McKee & Ostriker 2007). We will take an observational perspective, with emphasis on the interface between local and galactic scales. We focus on nearby galaxies and the MW, bringing in results from more distant galaxies only as they bear on the issues under discussion.

The remainder of this article is organized as follows. In the next subsection, we list some definitions and conventions that will be used throughout the paper. All observations in this subject rest on quantitative diagnostics of gas properties and star formation rates on various physical scales, and we review the current state of these diagnostics (§ 2, § 3). In § 4, we review those properties of Galactic star-forming regions which are most relevant for comparison to other galaxy-wide studies. In § 5, we review the star-forming properties of galaxies on the large scale, including the MW; much of this section is effectively an update of the more extended review presented in K98. Section 6 updates our knowledge of relations between star formation and gas (e.g., the Schmidt law) and local tests of these relations. The review concludes in § 7 with our attempted synthesis of what has been learned from the confluence of local and global studies and a look ahead to future prospects.

We list in § 1.3 some key questions in the field. Some have been or will be addressed by other reviews. Some are best answered by observations of other galaxies; some can only be addressed by observations of local star-forming regions in the MW. For each question we list sections in this review where we review the progress to date in answering them, but many remain largely unanswered and are referred to the last section of this review on future prospects.

Space limitations prevent us from citing even a fair fraction of the important papers in this subject, so we instead cite useful examples and refer the reader to the richer lists of papers that are cited there.

1.2. Definitions and Conventions

Here we define some terms and symbols that will be used throughout the review.

The term "cloud" refers to a structure in the interstellar medium (ISM) separated from its surroundings by the rapid change of some property, such as pressure, surface density, or chemical state. Clouds have complex structure, but theorists have identified two relevant structures: clumps are the birthplaces of clusters; cores are the birthplaces of individual or binary stars (e.g., McKee & Ostriker 2007). The observational equivalents have been discussed (e.g., Table 1 in Bergin & Tafalla 2007), and cores reviewed (di Francesco et al. 2007; Ward-Thompson et al. 2007a), and we discuss them further in § 2.2.

We will generally explicitly use "surface density" or "volume density", but symbols with Σ will refer to surface density, N will refer to column density, and symbols with n or ρ will refer to number or mass volume density, respectively. The surface density of HI gas is represented by ΣHI if He is not included and by Σ(atomic) if He is included. The surface density of molecular gas is symbolized by ΣH2, which generally does not include He, or Σmol, which generally does include He. For MW observations, He is almost always included in mass and density estimates, and Σmol may be determined by extinction or emission by dust (§ 2.3) or by observations of CO isotopologues (§ 2.4). When applied to an individual cloud or clump, we refer to it as Σ(cloud) or Σ(clump). For extragalactic work, ΣH2 or Σmol are almost always derived from CO observations, and inclusion of He is less universal. In the extragalactic context, Σmol is best interpreted as related to a filling factor of molecular clouds until Σmol ~ 100 M pc-2, where the area filling factor of molecular gas may become unity (§ 2.4). Above this point, the meaning may change substantially. The total gas surface density, with He, is Σgas = Σmol + µ ΣHI, where µ is the mean molecular weight per H atom; µ = 1.41 for MW abundances (Kauffmann et al. 2008), though a value of 1.36 is commonly used.

The conversion from CO intensity (usually in the J = 1 → 0 rotational transition) to column density of H2, not including He, is denoted X(CO), and the conversion from CO luminosity to mass, which usually includes the mass of He, is denoted αCO.

The term "dense gas" refers generally to gas above some threshold of surface density, as determined by extinction or dust continuum emission, or of volume density, as indicated by emission from molecular lines that trace densities higher than does CO. While not precisely defined, suggestions include criteria of a threshold surface density, such as Σmol > 125 M pc-2 (Goldsmith et al. 2008; Lada et al. 2010; Heiderman et al. 2010), a volume density criterion (typically n > 104 cm-3) (e.g., Lada 1992), and detection of a line from certain molecules, such as HCN. The surface and volume density criteria roughly agree in nearby clouds (Lada et al. 2012) but may not in other environments.

The star formation rate is symbolized by SFR or dot{M}*, often with units of M yr-1 or M Myr-1, and its meaning depends on how much averaging over time and space is involved. The surface density of star formation rate, Σ(SFR), has the same averaging issues. The efficiency of star formation is symbolized by є when it means M* / (M* + Mcloud), as it usually does in Galactic studies. For extragalactic studies, one usually means dot{M}* / Mgas, which we symbolize by є. The depletion time, tdep = 1 / є. Another time often used is the "dynamical" time, tdyn, which can refer to the free-fall time (tff), the crossing time (tcross), or the galaxy orbital time (torb). In recent years, it also has become common to compare galaxies in terms of their SFR per unit galaxy mass, or "specific star formation rate" (SSFR). The SSFR scales directly with the stellar birthrate parameter b, the ratio of the SFR today to the average past SFR over the age of the galaxy, and thus provides a useful means for characterizing the star formation history of a galaxy.

The far-infrared luminosity has a number of definitions, and consistency is important in converting them to dot{M}*. A commonly used definition integrates the dust emission over the wavelength range 3 - 1100 (Dale & Helou 2002), and following those authors we refer to this as the total-infrared or TIR luminosity. Note however that other definitions based on a narrower wavelength band or even single-band infrared measurements are often used in the literature.

We refer to mass functions with the following conventions. Mass functions are often approximated by power laws in logM:

Equation 1


In particular, the initial mass function (IMF) of stars is often expressed by equation 1 with a subscript to indicate stars (e.g., M*).

Structures in gas are usually characterized by distributions versus mass, rather than logarithmic mass:

Equation 2


for which α = γ + 1. Terminology in this area is wildly inconsistent. To add to the confusion, some references use cumulative distributions, which decrease the index in a power law by 1.

The luminosities and masses of galaxies are well represented by an exponentially truncated power-law function (Schechter 1976):

Equation 3


The parameter α denotes the slope of the power-law function at low luminosities, and L* represents the luminosity above which the number of galaxies declines sharply. For the relatively shallow slopes (α) which are typical of present-day galaxies, L* also coincides roughly with the peak contribution to the total light of the galaxy population. As discussed in § 5.2, the Schechter function also provides a good fit to the distribution function of total SFRs of galaxies, with the characteristic SFR in the exponential designated as SFR*, in analogy to L* above.

The term "starburst galaxy" was introduced by Weedman et al. (1981), but nowadays the term is applied to a diverse array of galaxy populations. The common property of the present-day populations of starbursts is a SFR out of equilibrium, much higher than the long-term average SFR of the system. No universally accepted quantitative definition exists, however. Some of the more commonly applied criteria are SFRs that cannot be sustained for longer than a small fraction (e.g., ≤ 10%) of the Hubble time, i.e., with gas consumption timescales of less than ≪1 Gyr, or galaxies with disk-averaged SFR surface densities Σ(SFR) ≥ 0.1 M yr-1 kpc-2 (§ 5.2). Throughout this review we will use the term "quiescent star-forming galaxies" merely to characterize the non-starbursting galaxy population. Note that these local criteria for identifying starbursts are not particularly useful for high-redshift galaxies; a young galaxy forming stars at a constant SFR might more resemble a present-day starburst galaxy than a normal galaxy today.

1.3. Questions

  1. How should we interpret observations of the main molecular diagnostic lines (e.g., CO, HCN) and millimeter-wave dust emission? How does this interpretation change as a function of metallicity, surface density, location within a galaxy, and star formation environment? (§ 2.4, § 5.1, § 7.3)
  2. How does the structure of the ISM, the structure of star-forming clouds, and the star formation itself change as a function of metallicity, surface density, location within a galaxy, and star formation environment? (§ 2.4, § 5, § 7.3).
  3. How do the mass spectra of molecular clouds and dense clumps in clouds vary between galaxies and within a galaxy? (§ 2.5, § 5.1, § 7.3)
  4. How constant is the IMF, and how are star formation rate measurements affected by possible changes in the IMF or by incomplete sampling of the IMF? (§ 2.5, § 3.3, § 4.2, § 6.4)
  5. What are the limits of applicability of current star formation rate tracers? How are current measurements biased by dust attenuation or the absence of dust, and how accurately can the effects of dust be removed? How do different tracers depend on metallicity, and what stellar mass ranges and timescales do they probe? (§ 3)
  6. How long do molecular clouds live and how can we best measure lifetimes? Do these lifetimes change systematically as functions of cloud mass, location in a galaxy, or some other parameter? (§ 4.3, § 7.3)
  7. Do local observations provide any evidence for bimodality in modes of star formation, for example distributed versus clustered, low-mass versus high-mass star formation? (§ 4, § 7.3)
  8. On the scale of molecular clouds, what are the star formation efficiencies [є = M* / (M* + Mcloud)], and star formation rates per unit mass (є = dot{M}* / Mcloud), and do these efficiencies vary systematically as functions of cloud mass or other parameters? (§ 4, § 7.3)
  9. How do spatial sampling and averaging affect the observed form of the star formation relations expressed in terms of total, molecular, or dense gas surface densities? How well do dot{M}* diagnostics used in extragalactic studies work on finer scales within a galaxy? (§ 3.9, § 6, § 6.4, § 7.3)
  10. Are there breakpoints or thresholds where either tracers or star formation change their character? (§ 2, § 3, § 7)

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