Introduction

Giant interstellar clouds are the most massive chemical ‘factories’ in our Galaxy containing around 80 molecules presently identified (neglecting isotopic variants) and ranging in complexity from H2 and CO to large saturated molecules such as ethanol, CH3CH2OH, and highly unsaturated cyanopolyynes including the 13-atom chain HC11N. Millimetre and sub-millimetre observations of interstellar molecules allow one to probe the densities, temperatures, and dynamics of interstellar clouds and can give information on the initial conditions for star formation. It is also of interest to understand the chemistry of these molecules since chemical kinetic modelling can be used together with observational data to constrain uncertain parameters such as elemental abundances and the cosmic ray ionisation rate. An understanding of deuterium fractionation in interstellar molecules can be used to determine the D/H ratio and thus has a bearing on cosmological models of the origin of the universe.

In recent years models of increasing chemical, physical, and computational complexity have been developed to study molecular formation and destruction in various astronomical regions. Models of interstellar cloud chemistry can be divided roughly into two classes: (1) steady-state models, in which chemical abundances are calculated through solving a coupled system of non-linear algebraic equations; and (2) time-dependent models, in which the variations of abundances as a function of time are followed through solving a coupled system of stiff, non-linear, first-order, ordinary differential equations. Traditionally, steady-state models have been the dominant tool for studying chemistry in diffuse clouds which, in the absence of shocks, reach steady-state well within their lifetimes.

Introduction

The first unambiguous detection of molecules in galaxies other than the Milky Way was that of OH in absorption toward the nuclear continuum sources in NGC 253 and M82 (Weliachew 1971). Subsequently, all of the species detected in the Milky Way that show strong emission or absorption lines have now been detected in other galaxies. Although extragalactic molecules have been observed in regions of the spectrum other than the radio, notably the vibrational–rotational lines of H2 in the infrared (Thompson, Lebofsky and Rieke 1978), observations remain largely the domain of radioastronomers. Furthermore, the largest fraction of all extragalactic molecular observations have been carried out using the J = 1−0 transition of CO, because it is the easiest transition to detect in most galaxies.

Initial progress in the field was slow, hampered by the simultaneous requirements of large bandwidth, high sensitivity and high angular resolution. The technical requirements were largely overcome in the early 1980s and the field has now burgeoned into one where developments are occurring so fast, on so many fronts that it is impossible in a short review to do justice to all of them. Accordingly, I have chosen a small subset of what I think are among the most interesting recent developments in the study of molecules in other galaxies with the aim of focussing on a few problems of wide astrophysical interest. A good introduction to the field is the comprehensive review by Morris and Rickard (1982).

Introduction

There is an increasing body of evidence for a population of large molecules in the interstellar medium. Large molecules have been suggested as the source of near infrared continuum radiation and the near infrared emission bands observed in reflection nebula, planetary nebula, HII regions and active galaxies (Duley and Williams 1981, Léger and Puget 1984, Allamandola, Tielens and Barker 1985, d'Hendecourt et al. 1986, Barker, Allamandola and Tielens 1987). Large molecules have also been proposed as the carriers of the diffuse interstellar bands (vander Zwet and Allamandola 1985, Léger and d'Hendecourt 1985, Crawford, Tielens and Allamandola 1985). The large molecules proposed have between 30 and 100 atoms with suggested abundances in the range 10−7−10−6 relative to hydrogen.

Mathis, Rumpl and Nordsieck (1977; MRN) determined a size distribution for grains that would fit the extinction measured to many sources. However, existing extinction measurements do not extend far enough into the ultraviolet to infer the population of the smallest grains. An extrapolation of the MRN distribution to very small grains having between 30 and 100 atoms gives an abundance roughly the same as that proposed for the large molecules, suggesting that the large molecules may be an extension of the interstellar grain distribution.

It appears that a population of small grains or large molecules exists in much of the interstellar medium. We will use the terms large molecules and small grains interchangeably to refer to clusters of between 30 and 100 atoms.

Introduction

The wide variety of ionized and neutral molecules detected in diffuse and dense interstellar clouds are mainly synthesized in the gas phase by many sequential and parallel positive ion–neutral reactions. The product positive ions may then be neutralized by proton transfer or charge transfer reactions with ambient molecules and/or via dissociative positive ion–electron recombination reactions and/or (as has recently been proposed) via mutual neutralization (positive ion–negative ion recombination) reactions involving large negatively-charged molecules. The basic gas phase ion chemistry was outlined in the early 1970s (Solomon and Klemperer 1972, Herbst and Klemperer 1973, Dalgarno 1975, Dalgarno and Black 1976) and since that time more sophisticated quantitative ion–chemical models have been developed to describe the routes to, and relative abundances of, the increasing number of complex molecules detected in interstellar clouds (Leung, Herbst and Huebner 1984, Millar and Nejad 1985, Dalgarno 1986, van Dishoeck and Black 1986, Millar, Leung and Herbst 1987). These models require the input of a large amount of kinetic data which (ideally) has been obtained at the low temperatures pertaining to interstellar clouds. In particular, the rate coefficients and products of many positive ion–neutral reactions and recombination reactions are required. Hence, a good deal of effort has been made to design and exploit laboratory experiments to provide these data (Huntress 1977, Smith and Adams 1981a, Anicich and Huntress 1986, Adams and Smith 1987a, Rowe 1988, Adams and Smith 1988a). This has led to a better understanding of the molecular processes that can occur under the extreme conditions of low pressure and low temperature in quiescent interstellar clouds.

Introduction

Fast shocks destroy molecules. Whereas the compressive and heating effects of slow shocks profoundly alter the chemical composition of a molecular gas but leave its molecular nature intact, interstellar shocks travelling faster than about 50 km s−1 result in the complete and very rapid dissociation of any pre-existing molecules by collisional processes (see Figure 3 in McKee, Chernoff and Hollenbach 1984), and if the shock velocity exceeds 70 km s−1, the atomic dissociation products are then largely ionised. Shocks faster than 80 km s−1 generate enough ultraviolet radiation to destroy molecules by photoionisation before they even reach the shock front (Hollenbach and McKee 1989).

Dissociative shocks may be present in the interstellar medium wherever gas is moving supersonically at velocities greater than 50 km s−1. Such velocities of bulk motion may result from the outflow of material from young stellar objects, from stellar winds, or from supernova explosions. Clearly, dissociative shocks are intimately associated with stars: they may be generated during the birth, throughout the life and by the explosive death of stars. Emission from fast shocks may serve as a tracer of active star formation in the Galaxy.

Despite the extremely destructive effects of fast shocks, molecules can reform behind a passing dissociative shock before the shocked gas has cooled to the temperature of the ambient, unshocked medium. Free electrons produced behind the shock front initiate the process by catalysing H2 formation. The rapid return of a dissociated shocked plasma to the molecular phase is a striking demonstration of the tendency toward molecule formation in the interstellar medium.

Introduction

Infrared line emission from vibrationally excited molecular hydrogen (H2) has been observed in many objects in the Galaxy including planetary and reflection nebulae and various molecular cloud complexes such as those in Orion and the vicinity of the Galactic center. Similar emission has also been detected in diverse extragalactic sources such as Seyfert and interacting galaxies and objects in the Small Magellanic Cloud. Most interstellar molecular hydrogen does not emit substantial vibrational emission since it exists in cold clouds at temperatures less than 100 K while the vibrational energy levels lie many thousands of degrees above the molecular ground state. Considerable infrared vibrational emission is produced, however, in regions which are heated to sufficiently high temperatures where the hydrogen molecules are thermally excited by collisional processes, or in regions where an efficient nonthermal molecular excitation mechanism is operating.

Many of the observed H2 emitting regions are physically associated with sources of intense ultraviolet radiation. Photons with wavelengths longward of the Lyman limit can escape the ionized clouds of hydrogen gas that usually surround the radiation sources and penetrate into neutral gas clouds called photodissociation regions. The thermal and chemical structures of these regions are critically influenced by the ultraviolet radiation. The molecular hydrogen that is present in these clouds is vibrationally excited by the discrete absorption of ultraviolet photons, and it may also be collisionally excited in warm gas heated by the ultraviolet radiation. The radiative decay of the excited molecular hydrogen produces an infrared spectrum that depends on a variety of physical parameters which may vary considerably from one interstellar photodissociation region to another.

Introduction

This chapter describes some aspects of the structure, motions, energies, and evolution of molecular clouds, from the viewpoint of observations and simple models. The emphasis is on observations of nearby clouds, and on physical, rather than chemical properties. The approach is partly pedagogical and partly a summary of recent results. Recent reviews on related subjects are those of Blitz (1987; diffuse molecular clouds), Scalo (1987; turbulence), Mouschovias (1987; magnetic effects), Shu, Adams and Lizano (1987; cloud physics and star formation), and Larson (1988; large-scale aspects of cloud and star formation). A review with a similar viewpoint to this one, but with older information, appears in Myers (1987).

We consider a molecular cloud to be a collection of interstellar gas and dust, whose gas has a substantial molecular component, and whose mean density inside an observationally definable boundary exceeds that outside the boundary. In this chapter we do not discuss galaxies, or molecular clouds in galaxies other than the Milky Way. Further information on these subjects is in Chapter 2.

The main constituents of molecular clouds, other than molecules, are stars, the ‘cores’ or condensations that form stars, dust grains, ions, and atoms. In this chapter we focus on the role of stars and cores.

The material in this chapter is organized into discussions of cloud structure and kinematics (Section 18.2), cloud energetics (Section 18.3), and time scales and evolution (Section 18.4).

Introduction

Observations of infrared line emission from molecular hydrogen in astronomical sources have gone from the novel to the commonplace in a time that is short relative to most timescales for the advancement of astrophysical knowledge. In the face of the current onslaught of observations of H2, it may seem surprising that only a dozen years ago, when the 2 µm lines were first detected, it is reported to have taken their discoverers several months to identify the emitting species. Excited H2 continues to be detected in new and surprising places. In the Galaxy, H2 line emission is found at interfaces between young stellar winds and the interstellar medium, where it first was discovered by Gautier et al. (1976), in reflection nebulae (Gatley et al. 1987), in supernova remnants (e.g., Burton et al. (1988)), even including the Crab Nebula (Graham, Wright, and Longmore 1989), in planetary and proto-planetary nebulae (Treffers et al. 1976), and in the nucleus (Gatley et al. 1984). Beyond the Galaxy H2 line emission is found in Seyfert, as well as starburst galaxies (Thompson, Lebofsky, and Rieke 1978, Joseph et al. 1986, Fischer et al. 1987), in individual HII regions of normal spiral galaxies (Israel et al. 1989), and in interacting and merging galaxies (e.g., Joseph et al. 1986). Closer to home, H2 line emission has been detected very recently in the aurora of Jupiter (Trafton et al. 1988). In all of these examples, because of the unique physical conditions which must be satisfied in order that its infrared line emission be detectable, H2 lines provide important information about environments that are difficult to study by other means.

Introduction

Ultraviolet radiation is a crucial ingredient in any theory of interstellar chemistry. In the interplay of molecule formation and destruction processes, ultraviolet photons adopt a multiple role, destroying neutral species on the one hand, while creating chemically reactive ions and depositing thermal energy on the other. It has long been recognized (e.g. Stief et al. (1972)), that dust in a cloud's outer layers attenuates ambient Galactic ultraviolet starlight, thereby enhancing the survival of molecules against photodestruction. Unfortunately, however, the degree of attenuation is sensitive to the grain scattering properties, which are not well determined at ultraviolet wavelengths (Sandell and Mattila 1975, Leung 1975, Whitworth 1975, Bernes and Sandqvist 1977, Sandell 1978, Flannery, Roberge, and Rybicki 1980). Since even a small amount of ultraviolet radiation has profound consequences in dark regions, the chemical and ionization balance of such regions has remained uncertain.

The early studies of dust shielding may have been overly pessimistic about uncertainties, however, as noted by Chlewicki and Greenberg (1984a,b). This is due in part to the existence of strong constraints on grain properties that follow from secure observational data, and also to the discovery of the chemical consequences of the ultraviolet emission associated with gas-cosmic ray interactions (Prasad and Tarafdar (1983); see also Chapter 16). The interaction produces an ultraviolet field in clouds which, at great depths, destroys molecules more rapidly than cosmic rays or attenuated starlight. As a result, the role of starlight is restricted to a relatively narrow region near a cloud's surface, where the effects of uncertainties in grain properties are moderate.

Introduction

Shock waves are ubiquitous in the interstellar medium (ISM) because efficient radiative cooling allows interstellar gas to cool to temperatures low enough that the sound speed is small compared to the velocities of disturbances in the ISM, such as cloud–cloud collisions, bipolar outflows, expanding HII regions, and supernova explosions. Shock waves in dense molecular gas are almost always radiative: The relative kinetic energy of the shocked and unshocked gas is converted into radiation, and since the radiating gas is dense, it is very bright. Because much of the mass in molecular clouds is obscured by dust, the emission from shocks provides a powerful probe of energetic activity occurring in these clouds. In particular, stars inject large amounts of energy into their surroundings in the process of formation, giving rise to bipolar outflows with velocities in excess of 100 km s−1, characteristic of stellar escape velocities (Lada 1985). Intense maser emission in the 1.35 cm line of water is also observed to be associated with newly formed stars, particularly massive stars, with velocities of tens to hundreds of kilometers per second (Genzel 1986). Understanding the structure and spectrum of the shocks associated with these high velocity flows in dense molecular gas is thus a prerequisite for unraveling the complex physical processes attending the birth of stars.

Early studies of shocks in molecular clouds assumed that the neutrals and ions were tied together into a single fluid, and that the shock front was an abrupt transition on the scale of the molecular mean free path (e.g., Field et al. (1968), Hollenbach and McKee (1979)).

Introduction

Interstellar chemistry began to be studied in a fairly serious way when, in the late 1960s and early 1970s, it was demonstrated that a wide variety of molecules existed in dense molecular clouds. At first the main effort was in identifying the main chemical routes by which molecules were formed and destroyed. It was realized that, even in dark molecular clouds where starlight is excluded, cosmic rays may penetrate and cause ionizations which drive a chemistry which would otherwise ‘run down’. This chemistry would, therefore, be largely one of positiveions and molecules. This early recognition met with great success and – although the level of ionization in molecular clouds remains uncertain – the detection of interstellar ions such as HCO+ and N2H+ is strong support for positive ionneutral molecule chemistry. Models of interstellar chemistry involving hundreds or even thousands of reactions are now routinely studied: some of these reactions may be important.

These early studies, understandably, concentrated on the chemistry. They deliberately made the dynamics as simple as possible. Thus, uniform density and temperature were usually invoked, in geometrically convenient shapes such as semi-infinite slabs, or spheres. Steady-state calculations were often performed, without a full consideration of the applicability of steady-state. Later studies showed that it might take around 30 million years to achieve steady-state in molecular clouds, and it was realised that such extended periods might not be available in interstellar clouds.

At the same time, detailed observations were indicating that molecular clouds were far from simple objects.

Historical perspective

The study of interstellar chemistry started, appropriately, about 60 years ago. In 1926, Eddington discussed in his remarkable Bakerian Lecture the possibility of molecule formation and absorption in dark nebulae. At that time, only atomic species had been identified in interstellar space through their narrow absorption lines superposed on the spectra of background stars. In the next decade, several new interstellar features were detected which could indeed be ascribed to molecules: CH, CH+ and CN.

In spite of this early success, no other molecule was found in the interstellar gas for the next 25 years, until OH was detected in 1963 by its radio emission lines. In the next two decades, more than 70 different interstellar molecules were identified by centimeter and millimeter wavelength techniques. However, these radio emission line studies were mostly concerned with dense and dark clouds, whereas the early absorption line observations probed much more diffuse gas.

Although a wide variety of interstellar molecules has now been detected in dark clouds, still only a handful of molecules has been found in diffuse clouds. The molecules H2, HD, OH and CO were discovered in the 1970s by their absorption lines in the ultraviolet through rocket experiments and by the Copernicus satellite. Since the detection of C2 in 1977 by ground-based techniques, however, no new molecule has firmly been identified in diffuse clouds. The list of molecules sought but not detected is considerably longer and includes such interesting species as NH, HC1, NaH, MgH, H2O and C3 (see van Dishoeck and Black (1988a) for a recent summary).

Introduction

Over the past ten years, observations of the far-infrared and submillimeter fine structure emission from neutral oxygen and carbon, and singly ionized carbon and silicon, have revealed the presence of a dense (nH ≤ 103 cm−3), intermediate temperature (60 ≥ Tgas ≥ 1500 K) component of the interstellar medium in which a substantial fraction of the gas is atomic (see Table 14.1). In addition, observations of atomic fine structure emission from regions which are mostly molecular have helped to define the nature of the shocks which frequently accompany star formation and to challenge long-held beliefs about the structure and chemistry of molecular clouds.

Emission from atomic and ionic species with ionization potentials less than that of hydrogen arises predominantly in diffuse HI clouds, the photodissociated surfaces of molecular clouds, and the warm gas downstream of passing shock waves, generated deep within molecular clouds by the outflows from newly formed stars. The latter two regions, of primary interest in this contribution, are characterized by gas temperatures of 60–2000 K. Since H and He have no lowlying levels which can effectively cool the gas in this temperature range, the fine structure transitions of OI (63, 146 µm), C+ (158 µm), and Si+ (35 µm) generally dominate the cooling and thus serve as important diagnostics of these regions. An indication of the cooling power in just the [OI] 63 µm and [CII] 158 µm lines is given in Table 14.2; between 0.1 and 1% of the total luminosity of these sources escapes in these two lines.

Introduction

The high observed column densities of CH+, one of the first identified (Douglas and Herzberg 1941) interstellar molecules, and of CO apparently indicate that existing static, equilibrium models do not provide adequate descriptions of the natures of diffuse molecular interstellar clouds. (See Chapter 3.) It has been argued that velocity structures in lines formed in such clouds provide evidence for the existence of shocks in them (e.g. Crutcher (1979), but see the detailed assessment by Langer in Chapter 4). If such shocks do exist, they will drive the production of detectable column densities of a number of chemical species.

The chemistry in shocked gas can be exceptionally rich since many reactions which, because they are endothermic or have activation barriers, are unimportant in cool, static gas, can proceed in shocked gas. For instance, the endothermic reactions C+ + H2 → CH+ + H (Elitzur and Watson 1978a) and S+ + H2 → SH+ + H (Millar et al. 1986) can initiate hydrogen abstraction sequences in shocked gas but are unimportant in static, cool diffuse clouds. A neutral–neutral sequence (Aannestad 1973) which is of no relevance to low temperature chemistry but which plays a major role in shock chemistry is O + H2 → OH + H; OH + H2 → H2O + H. The fractional abundances of CH+ and OH are high in some diffuse cloud shocks, and SH+ may serve as a diagnostic of shocks.

Collisionally induced rotational excitation of molecular hydrogen can also occur in diffuse cloud shocks.