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.

Introduction

In 1973 Drake and Sagan proposed a SETI frequency standard of V0 ∼ 56 GHz tied to the observed cosmic microwave background hv0 = kT0, where T0 is the current temperature of the cosmic microwave background. They noted that a transmitting civilization in a distant galaxy will, however, have transmitted its signal to the Earth at an earlier cosmological epoch when T was larger than is measured today, tending to increase the ‘natural’ frequency, but that the cosmological Doppler effect will tend to decrease the frequency. Not knowing of their work, I proposed this same frequency standard (hv0 = kT0: Gott, 1982) in the first edition of this book. I had noticed that the two effects mentioned above in fact cancel each other out exactly (which was a new result), so that this frequency standard was indeed universal. If a transmitting civilization is at a redshift z, it will observe a microwave background temperature of T1 = T0 (1 + z) and will emit signals at a frequency of hve = kT1 = kT0 (1 + z), but because of the cosmological Doppler shift we will observe these transmitted photons at a frequency hv0 = hve (1 + z)–1 so that we observe hv0 = kT0 (1 + z) (1 + z)–1 = kT0 ∼ 56 GHz independent of the redshift of the emitting civilization.

The purpose of a second edition of Where Are They? is to enlarge upon and update issues that were debated in 1978 at a two-day Symposium of the same name. As might be expected, comparison of the present book with the first edition shows that relatively little has changed in the field of interstellar travel and colonization – there are not many interstellar travelers among us. By way of contrast, because the world is full of biologists and astronomers, there have been many new experiments and insights in these fields. We are especially pleased that three distinguished biologists – Drs Diamond, Joyce and Mayr – have contributed new chapters to the present volume. Typically, biologists appear to be less sanguine about the likelihood of abundant intelligent life in the universe than are engineers and physicists.

At the time of publication of the first edition, the only things that humans knew of with certainty that orbited stars other than the Sun were other stars. Then, in 1983, NASA's IRAS satellite discovered that many nearby stars similar to our Sun emit infrared (heat) radiation well in excess of that expected from their visible surfaces. In a chapter in the first edition entitled ‘Searches for electromagnetic signals from extraterrestrial beings’ I discussed the possibility that IRAS might discover a so-called ‘Dyson Sphere’.

In one school of thought it is customary to begin discussions of galactic life by appeal to Drake's equation and then to proceed to a detailed examination of the numerical magnitude of one or more of the string of factors whose values have to be estimated. An example of this procedure is furnished by Michael H. Hart's analysis (Chapter 22), in which he concentrates on the probability that 600 or more nucleotides might line up in the right order; then he proposes that one of the factors may be very much less than 10–30. Of course, 10–30 is already very small, and, if included as a factor in almost any expression having to do with the physical universe, will cut the product down to negligible size. In this application the conclusion is that the number of technological civilizations independently arising in a galaxy is very much less than 1. Well, this may be an excess of zeal, and many of those addicted to the use of Drake's equation would, in similar circumstances, have arranged for the product to emerge with an order of magnitude around unity because, after all, a calculation condemns itself if it seriously contradicts the possibility of the one technological civilization we know about, namely our own.

But is Drake's equation correct? It seems that it suffers from oversimplification - surely at least one plus sign ought to be there.

The next time you're outdoors on a clear night and away from city lights, look up at the sky and get a sense of its myriads of stars. Train your binoculars on the Milky Way and appreciate how many more stars escaped your naked eye. Then look at a photograph of the Andromeda nebula as seen through a powerful telescope to realize the enormous number of stars that escaped your binoculars as well. When all those numbers have sunk in, you're ready to ask: How many civilizations of intelligent beings like ourselves must be out there, looking back at us? How long before we are in communication with them, before we visit them or before we are visited?

Many scientists have tried to calculate the odds. Their efforts have spawned a whole new field of science termed exobiology – the sole scientific field whose subject matter has not yet been shown to exist. Since a summary of the calculations fills seven pages of the Encyclopaedia Britannica, what more could we learn by further speculation? I'll suggest, nevertheless, that woodpeckers offer a fresh perspective.

Exobiologists find the numbers in their subject matter encouraging. Billions of galaxies each have billions of stars. Many stars probably have one or more planets, and many of those planets probably have an environment suitable for life. Where suitable conditions exist, life will probably evolve eventually.

Introduction

The possibility that life, primitive or advanced, might exist in other parts of the universe has occupied the thoughts of scientists and laymen for thousands of years. One of the earliest was the statement by the ancient Greek philosopher Metrodorus of Chios around 400 b.c., who wrote in his book On Nature that: ‘It is unnatural in a large field to have only one shaft of wheat, and in the infinite Universe only one living world.’

In a.d. 1690 the famous Dutch physicist Christian Huygens wrote in his book Cosmotheoros that: ‘Barren planets, deprived of living creatures that speak most eloquently of their Divine Architect, are unreasonable, wasteful and uncharacteristic of God, who has a purpose for everything.’

In the nineteenth century, several proposals were made by different distinguished scientists. The most famous was mathematician Carl Friedrich Gauss, who proposed to establish contacts with advanced civilizations on other planets of our solar system, by planting a rectangular triangle with wheat in Siberia, with squares of pine trees at its three sides, to show that the Earth has intelligent beings that know the Pythagorean Theorem. None of these proposals, however, was implemented.

The modern era of the Search for Extra-Terrestrial Intelligence (SETI) started in 1959 with a paper to Nature by Cocconi and Morrison, which was followed soon after in the spring of 1960 by the first radio search by Frank Drake (Project OZMA), using the then new 85 foot radio telescope at the National Radio Astronomy Observatory in West Virginia.

Interstellar Travel and Extraterrestrial Intelligence

The success of several proposals to search for extraterrestrial intelligence (ETI) in the Galaxy (Cocconi & Morrison, 1959; Oliver & Billingham, 1971; Michaud, 1979) requires the existence of a large number of technologically competent cultures over a long period of time. For example, to expect to find one ETI within 1000 light-years in a perfectly efficient search would require about a million ETI in the Galaxy, each signalling for a million years. (Or it would require 108 ETI signalling 104 years, or 104 ETI signalling 108 years, etc.) Many people have asked why some of these ETI should not have taken advantage of their prolonged technological capability to find a method for interstellar travel and settlement of nearby stellar systems (see, e.g., Hart, 1975; Jones, 1976; Winterberg, 1979). If the initial problem of interstellar travel and settlement were solved, then it should become progressively easier for daughter settlements to eventually continue the process until every available stellar system in the Galaxy (including possibly our own) were inhabited.

The chances of this happening have been discussed extensively, often with minimal thought given to the physical requirements for interstellar settlement. In particular, it has been argued that interstellar settlement is either impossible (see, e.g., Purcell, 1960; Marx, 1973) or absurdly expensive (e.g. requiring trillions of man-years of effort to amass the nuclear fuel needed).