This chapter will focus on those aspects of polarimetric instruments that are peculiar to certain wavelength regions. The concepts discussed in previous chapters will be used freely. Non-polarimetric wavelength-peculiar concepts will generally be taken for granted, but a few are essential and must be recapitulated briefly.

Optical/infrared systems

Optical polarimetric instrumentation has a long history of development. Early polarimeters had errors at the level of a few tenths of a per cent at best, and polarization signals were small, so that polarimetry was very much a specialist craft. B. Lyot was the first to obtain very high accuracy by devising a modulator and using it on the Sun. For stars, the signals were generally so small that photon shot noise was appreciable, and there was little incentive to design sophisticated systems of unavoidably smaller throughput.

The situation has changed drastically within the last decade or two. Larger telescopes are available, CCD detectors now offer thousands of parallel channels of potentially very good accuracy, and improved modulators of high transmission have been devised. The higher signal levels have meant that greater resolution (spectral, temporal, spatial) can be used, and this has had the effect of increasing the degree of polarization provided by nature (less smearing of polarizations from neighbouring resolution elements); the end result is that (i) many more situations within astronomy can usefully be tackled by polarimetry without exceptional cost in telescope time and (ii) ‘common-user’ polarimetry is becoming available in the optical/near-infrared wavelength region (the ‘CCD domain’).

This chapter introduces the tools used by astronomers and instrument designers in describing the action of a medium on the polarization of the radiation passing through it. In the majority of situations encountered in astronomy, the phase of the wave is unimportant, and we need a way to describe the transformation of Stokes parameters, i.e. the changing polarization forms which support the flow of radiant energy. For cases where the phase of the polarized radiation is important (e.g. polarization effects within an optical interferometer, the focusing of a plane wave by a radio telescope, or the amplification of polarized radiation within an astronomical maser), an alternative formulation will be introduced (in section 4.2) that describes the transformation (including phase) of the electric field vibrations of two orthogonal 100% polarized waves (usually linear polarization). In this formulation, partial polarization cannot be handled, and we must make separate calculations for two orthogonal polarizations of the incident radiation, constructing the incoherent sum at the end. Shurcliff (1962, sections 8.6, 8.7, 8.9) details the early history of these two calculi and compares their fields of use; a concise statement of the relationship between the two calculi may be found in Stenflo (1994, section 2.6).

Mueller matrices

As discussed in chapter 2, the four Stokes parameters denote the flow of radiant energy in specific vibrations of the electromagnetic field, and all four are expressed in the same units.

An overview

The classical spectroscopic signature of mass loss, first reviewed in the literature by Beals (1950), is the so-called P Cygni line profile. This label has come to be attached to the profile shape in which blueshifted absorption sits alongside redshifted emission. In truth, the practice of describing just this configuration as ‘P Cygni’ does little justice to the rich variety of profile forms that are to be found in this famous star's optical spectrum—Beals himself put the case for 4 different profile types characteristic of ‘P Cygni stars.’ Interestingly, from the perspective of this collection of papers on line emission, these other forgotten types include forms that emphasise emission rather than absorption. Indeed, those of us who have taken spectra of P Cygni itself are painfully aware of just how strong the strongest emission features (in Hα, He i λ5876) really are!

Introduction

The history of ultraviolet (UV) spectroscopy in astronomy spans over three decades now and such observations have led to many discoveries regarding the physical nature of the entire gambit of astronomical objects. Hot astrophysical plasmas have line and continuum emission and absorption processes for which UV spectroscopy can probe the more energetic physical processes that cannot be studied adequately in the optical or infrared. In addition, studies of the UV spectral properties of cooler bodies, such as planetary atmospheres, comets, and interstellar dust provide important information on their physical state and composition.

This article concentrates on reviewing some of the techniques and results from the study of emission lines in astronomical UV spectroscopy. Given that the range of astronomical objects from the Earth's geocorona to quasars show UV emission lines and that during the past three decades over two thousand papers have appeared in the literature, including numerous conferences and books, a comprehensive review is unpractical.

Apart from stars and those objects which radiate reflected starlight, most of the objects in the Universe radiate an emission spectrum. It was the astronomers interest in analyzing the spectrum of the sun and other stars in the last century that motivated the development of radiative transfer, and with the newly formulated macroscopic relations of LTE early in this century, that led to our understanding of absorption spectra. The original observational stimulus for this activity had been Fraunhofer's study of the solar spectrum almost a century before.

Interest in emission-line spectra came later, when spectrographs coupled to telescopes enabled the spectra of fainter gaseous emission regions to be observed. They revealed a totally different type of spectrum than that which had been observed from stars. The fact that local thermodynamic equilibrium does not hold for emission regions has complicated the interpretation of their spectra. Huggins' initial discovery of ‘nebulium’ in gaseous nebulae and its subsequent identification with ionized oxygen by Bowen had demonstrated that rarefied conditions must pertain in nebulae. Stromgren's subsequent 1939 paper in the Astrophysical Journal was a landmark in demonstrating how far-UV continuum radiation from a hot star was absorbed by surrounding gas and converted into visible Balmer line radiation.

In the decades that followed, the realization that many interesting objects such as supernova remnants, active galactic nuclei, and quasars radiated an emission-line spectrum, motivated the analysis of emission regions.

Introduction

Gamma ray lines are the signatures of nuclear and other high energy processes occurring in a wide variety of astrophysical sites, ranging from solar flares and the interstellar medium to accreting black holes and supernova explosions. Their measurement and study provide direct, and often unique, information on many important problems in astrophysics, including particle acceleration, star formation, nucleosynthesis and the physics of compact objects.

The physical processes that produce astrophysical gamma ray emission lines are nuclear deexcitation, positron annihilation and neutron capture. Excited nuclear levels can be populated by the decay of long-lived radioactive nuclei as well as directly in interactions of accelerated particles with ambient gas. Nuclear deexcitation lines following radioactive decay have been seen from supernova 1987A (Matz et al. 1988; Tueller et al. 1990; Kurfess et al. 1992), from the supernova remnants Cas A (Iyudin et al. 1994) and Vela (Diehl et al. 1995), and the interstellar medium (Mahoney et al. 1984; Share et al. 1985; Diehl et al. 1994; 1995).

Introduction

Numerical simulations of emission line regions, whether photo or shock ionized, are a vital tool in the analysis and interpretation of spectroscopic observations. Models can determine characteristics of the central source of ionizing radiation, the composition and conditions within the emitting gas, or, for shocks, the shock velocity. Osterbrock (1989) and Draine & McKee (1993) review the basic physical processes in these environments.

Although numerical simulations are a powerful tool, this capability is somewhat mitigated by the complexity of the calculations. There will always be underlying questions regarding the astronomical environment (i.e., the shape of the ionizing continuum, inhomogeneities, or the composition of the gas) and uncertainties introduced by the evolving atomic/molecular data base. On top of this, however, the numerical approximations, assumptions, and the complexity of the simulations themselves introduce an uncertainty that cannot be judged from a single calculation.

With these questions in mind Daniel Péquignot held a meeting on model nebulae in Meudon, France, in 1985. This provided a forum where investigators could carefully compare model predictions and identify methods, assumptions, or atomic data which led to significant differences in results.

Introduction

In a pioneering study on the electron impact excitation of atomic oxygen, Seaton (1953) formulated the now well known close-coupling approximation of atomic collision theory, which he termed the “continuum state Hartree-Fock method”, reflecting the physical picture that the new method was an extension of the bound state method to the continuum region that encompassed electron-ion scattering and photoionization phenomena. For nearly three decades, the close coupling approximation has been widely employed to calculate the most accurate low-energy cross sections for excitation and photoionization, and radiative transition probabilities. Large computational packages were developed, mainly at University College London and the Queen's University of Belfast, to carry out the enormous task of fulfilling the needs of astrophysicists and plasma physicists. In particular, the R-matrix method developed by Burke and associates (Burke et al. 1971) has proved to be computationally very efficient for large-scale calculations.

A huge amount of radiative atomic data was produced, during last 10 years or so, under the auspices of an international collaboration of atomic physicists and astrophysicists, called the Opacity Project, led by Seaton (Seaton et al. 1994).

In the first half of this century many emission lines were or had been identified. Noteworthy moments were the identification of the Nebulium lines (λ 4959/5007) as forbidden lines of O++ (Bowen 1927) and of the strong solar green coronal line λ 5303 as due to Fe13+ (Edlen 1942). In addition, a first quantitative understanding of some aspects of nebular spectra was obtained: the Balmer decrement was calculated by Menzel and associates (1937), the temperatures of the central stars of planetary nebulae were inferred by Zanstra (1927), and the first information on elemental abundances in nebulae was gained.

In the second half of this century a much more detailed understanding of emission spectra was acquired. Emission lines assumed a fundamental role for the diagnostics of conditions in nebulae. As a result electron densities Ne and temperatures Te as well as elemental abundances became known in many objects. Excitation and ionization conditions in nebulae were found to be frequently radiative (photoionization), but shocks and perhaps fast particles were also found to play a role. Non-equilibrium conditions were seen to be important especially in the hot, tenuous plasmas revealed by X-ray observations: the ionization state was often different from that expected from the temperature, and even Te and the temperature of the proton gas could be different.

Chemistry was found to play a role in many emission nebulae. Numerous new molecules were observed, especially by radio observations in cool, dense media.

Introduction and overview of active galactic nuclei

Observations of emission lines in photoionized nebulae provide important diagnostics of the line emitting gas in three different ways: Line intensities are used to derive the physical conditions in the gas. Density, temperature, optical depth etc. are all related to emission line ratios and absolute fluxes. Line profiles are used to investigate the gas dynamics and the velocity field. Finally, line variability, when correlated with flux variations of the photoionizing continuum, are used to measure the gas distribution and the size of the emission line region.

Active Galactic Nuclei (AGN) are situated in the center of otherwise normal galaxies and show strong emission lines superimposed on strong nonstellar continua.

Highlights

For the reader who has only a couple of moments to spare, the strongest overall impressions from “analysis of emission lines” were (1) infrared and ultraviolet astronomy have merged with optical astronomy in their techniques and power, and no longer need to be considered separately (except that Dufour and Dinerstein do these things so well); (2) limited wavelength resolution keeps this from being the case yet in X-ray astronomy, though planned missions promise improvements (Mushotzky), while gamma ray emission, coming largely from nuclear rather than atomic processes, will continue to require very different approaches (Ramaty); (3) the enormous growth of detailed atomic data (Pradhan) and sophisticated techniques for handling the partial redistribution of photons across line profiles and other non-linear processes in radiative transfer (Hummer) means that current computing power is not yet able to implement the best calculations that we, in principle, know how to do, especially for intrinsically complex systems like supernovae (Pinto) and lumpy stellar winds (Drew); and (4) there is something reassuring about encountering a large body of astronomical endeavor to which it matters hardly at all whether or not the early universe was dominated by a Gaussian, Harrison-Zeldovich spectrum of adiabatic fluctuations in biased Cold Dark Matter.

Introduction

The theory of radiative transfer has made spectacular advances in the past decade, both in the understanding of fundamentals and in computational techniques. However, apart from the solar/stellar community, these important tools for the interpretation of astrophysical spectra are neither recognized nor effectively used. It is hoped that this brief overview will be useful in communicating the state of understanding and guiding potential users to the appropriate literature. This paper is not intended as a review, but as a discussion of two important developments related to Osterbrock (1962).

The role of radiative transfer theory in the quantitative interpretation of spectra seems not to be widely understood. The crucial importance of radiative transfer processes as the link between an astronomical object and the determination of its physical properties is discussed in Sect. 2.

Although the necessity of treating radiation scattered in spectral lines as non-coherent, i.e., experiencing slight shifts in frequency in each scattering, is well understood, the conditions under which one can employ the simplifying assumption of complete redistribution are less well known. This issue is discussed in Sect. 3, starting from the discussion in Osterbrock (1962). Sect. 4 contains a detailed comparison of numerical solutions of the transfer equation with various assumptions concerning treatment of redistribution.

The solution of the combined radiative transfer and statistical equilibrium equations for atomic models with a large number of levels, and in various geometrical configurations, lies at the heart of the quantitative astrophysical spectroscopy.

Introduction

From the study of the emission lines produced in galactic and extragalactic gaseous nebulae it has been possible to derive abundances of H, He, C, N, O, Ne, S and Ar. The chemical composition of these gaseous nebulae is needed to understand their physical conditions as well as their evolution. These abundances are also paramount to constrain evolutionary models of stars, galaxies and the universe.

Reviews and textbooks on the physical processes taking place in ionized nebulae have been presented by many astronomers, classic ones are those by Seaton (1960), Aller (1984) and Osterbrock (1989).

Some abundances have been determined based on detailed photoionization models while most abundances have been determined based on simple empirical methods. The input of a photoionization model consists of: a) a stellar radiation field, b) an electron density distribution, Ne(r), (which defines the geometry of the nebula, c) a dust distribution, Nd(r), and d) abundance distributions, which in most cases have been assumed homogeneous. The output consists of: a) a set of line intensities, b) the electron temperature distribution, Te(r), and c) the ionization structure.

Introduction

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 ≤ 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.

Introduction

Molecules have been detected in a broad range of astronomical objects—the atmospheres of stars, diffuse, translucent and dense interstellar clouds, photon-dominated regions (PDRs), circumstellar shells, HII regions, planetary nebulae, stellar outflows, stellar winds, jets, Herbig-Haro objects, novae, supernova remnants and the eject a of Supernova 1987a. Their presence is a controlling force in the determination of the thermal balance, the ionization structure, the dynamics and the evolution of the entities in which they reside. Molecules provide unique diagnostic probes of the physical nature of their environment, yielding information on the densities, the temperatures, the magnetic fields, the velocities, the isotopic composition, the radiation fields, the masses and the ages.

In dense molecular clouds where star formation takes place, molecules have been detected in remarkable diversity. Over twenty of them have been detected in external galaxies. Many additional species have been discovered in circumstellar shells.

I will not attempt to survey the myriad ways in which molecules serve as diagnostic probes. I will instead make an arbitrary personal selection, beginning not with emission lines but with the absorption lines of CN which provided the first measurement of the temperature of the cosmic blackbody background radiation field.

Absorption by CN

Absorption by the CN molecule has been measured towards many stars. Several lines have been observed showing that CN is present in its low-lying rotational levels with rotational quantum numbers j = 0, 1 and 2.

Introduction

Let me preface this paper to say how honoured I am to have been given this opportunity to pay tribute simultaneously to two of the principal sources of scientific inspiration of my career. In my attempts over the past twenty years to understand and to analyse the optical and UV spectra of shock-excited plasmas, Don's books (1974, 1988) have been invaluable to both myself and to my students. In Australia we used to refer to the Physics of Gaseous Nebulae somewhat irreverently as “the new testament” to distinguish it from the earlier work by Aller! The famous diagnostic diagrams of Baldwin, Phillips & Terlevich (1981), of Veilleux & Osterbrock (1987), and of Osterbrock, Tran & Veilleux (1992) provide both an inspiration, and a powerful means of distinguishing between various excitation mechanisms.