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The phrase "infrared emission lines," like so many other topics covered at this meeting, is extremely broad. Therefore I will begin by defining and limiting the scope of this review. I will discuss only lines that arise from gas-phase atoms, ions, and molecules, and will not include spectral features produced by interstellar dust. In this review, "infrared" will mean the spectral region 1-200 µm, which corresponds to certain types of astronomical detectors; shorter wavelengths will be called "far red" (with apologies to Don Osterbrock), and longer wavelengths, "submillimeter." Another boundary condition is that I will restrict myself to low densities n leq 108 cm-3. Apart from these constraints, I will attempt to be as general as possible, though with no pretentions to completeness. This review is organized by physical properties and methods of analysis, rather than by class of astronomical object. In selecting applications I have tried to sample diverse areas of astronomy, in order to illustrate the breadth and power of the various methods for measuring and constraining astrophysically interesting parameters.

One might begin by asking "why infrared?" Perhaps it is not quite as necessary to raise this question as it was a decade or two ago, when infrared instrumentation was relatively primitive, and one needed to justify the extra effort required to make observations in the infrared region. However, it remains true that infrared emission lines have a number of special characteristics that translate into distinctive advantages, and a few limitations. Some special qualities of infrared emission lines include: (1) access to regions with high visual extinctions; (2) avoidance of the temperature sensitivity that is essentially intrinsic to most optical and ultraviolet emission (hereafter, optical/uv) lines; and (3) and the ability to sample species that lack observable lines in the optical/uv, either because of the vagaries of atomic structure or because the gas is too cool to emit in lines from high energy levels.

Most of the observed infrared lines (other than recombination lines of H and He) connect fine-structure levels within ground state pn configurations of abundant ions (where n is an integer between 1 and 5). A good synopsis of the physics of fine-structure lines can be found in Simpson (1975), an influential paper in the 1970's, a period of rapid advances in infrared spectroscopy. Other discussions of these lines can be found in Aller (1984, ch. 5), Watson (1985), and Osterbrock (1989, chs. 3 and 5). What follows here are just a few general points.

The upper levels of the fine-structure lines are populated by collisions, usually with electrons. For upper levels with small excitation energies chi, the populations are substantial. Whereas the excitation energies of the optical/uv lines are larger than the typical energy of the free electrons, for the infrared lines chi = hnu << kT and the exponential term is nearly equal to unity. The ratio of populations of an excited level (1) to the ground level (0) in a two-level atom is

Equation 1.1   (1.1)

(after Osterbrock 1989, eq. 3.25), where A10is the radiative transition probability, q01 and q10 are the collisional excitation and de-excitation probabilities respectively, and Omega(T) is the "collision strength," which depends only weakly on temperature. For lines in which both levels belong to the same electronic configuration, downward radiative transitions are forbidden by electric dipole selection rules. Combined with their low frequencies, this yields low transition probabilities, A10 leq 10-2 s-1, for the fine-structure lines. However, the large populations of these levels tend to compensate for the low A values, making the net volume emissivities of the infrared lines (in energy units) comparable to those of the optical/uv lines, which arise from much smaller populations in higher-energy levels.

From equation 1.1 one can see that the small excitation energies of the infrared lines insure a weak dependence of the emergent line strengths on electron temperature Te compared to the optical/uv lines. This has both advantages and disadvantages: it is helpful for deducing ionic abundances, since one need not know the exact value of Te; however, such lines therefore give no information about the temperature. The low transition probabilities, combined with high collisional de-excitation rates, lead to small critical densities, where, for level i,

Equation 1.2   (1.2)

(Osterbrock, eq. 3.31). As a result, the infrared lines are sensitive to density and clumping (i.e., Rubin 1989). They are good density indicators for low-density regions, but "saturate" to asymptotic values at relatively low densities, whereupon they become good mass tracers.

The history of the observations of the infrared emission lines more or less begins in the 1970's, when these lines were first observed in bright, compact H II regions and planetary nebulae. By 1982, the list of detected lines ran to fifteen (e.g., Dinerstein 1983). In the 1980's infrared lines began to be observed from a more diverse group of sources, such as novae, supernovae and supernova remnants, and photodissociation regions. Within the last few years, there has been a rapid growth in infrared spectroscopy of extragalactic sources, including active galactic nuclei and starburst galaxies. The fact that it is no longer feasible or sensible to present a comprehensive list of detected infrared emission lines in a review such as this one is a measure of the maturity of the field and the diversity of the astronomical sources and regions that benefit from being studied with infrared line spectroscopy.

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