The limitations of photometric techniques

The photometric all-sky surveys conducted in the EUV by the ROSAT WFC and EUVE have been sources of important information concerning the general properties of groups of objects contained in the EUV source population, including late-type stars, white dwarfs, cataclysmic variables and active galactic nuclei. However, when considering individual objects in detail, the amount of information that can be extracted from three or four such data points is limited. For example, if heavy elements are present in the atmosphere of a hot white dwarf, the survey data can only give an indication of the level of opacity and are unable to distinguish between the possible species responsible and especially whether or not helium is present. Similarly, in studying the emission from stellar coronae, only crude estimates of conditions in the plasma can be made and usually only when simplifying assumptions such as the existence of a single temperature component are incorporated into the analysis. The overwhelming advantage of spectroscopic observations lies in the ability to study individual spectral features or blends of features, giving a more detailed picture of the underlying physical processes responsible for the EUV emission.

The Extreme Ultraviolet Explorer spectrometer

The main components of EUVE have been described in detail in chapter 4 with the exception of the spectrometer. This instrument made use of part of the converging beam of the Wolter type II deep survey telescope, intercepting this with three variable line space reflection gratings (Hettrick et al. 1985).

The Extreme Ultraviolet Explorer

S. Bowyer and his team at the University of California, Berkeley, the pioneers of EUV astronomy, originally proposed the Extreme Ultraviolet Explorer (EUVE) mission to NASA in 1975. Selected for development in 1976, it was eventually launched in June 1992, almost exactly two years after ROSAT. The science payload was developed and built by the Space Sciences Laboratory and Center for EUV Astrophysics of the University of California at Berkeley. Like the WFC, a principal aim of EUVE was to survey the sky at EUV wavelengths and to produce a catalogue of sources. However, the two missions differed in several respects. While the WFC had a filter complement allowing it to observe at the longer wavelengths of the EUV band (P1 and P2 filters), the survey was only conducted at the shorter wavelengths from 60 Å to ≈200 Å. By contrast, the EUVE survey was carried out in four separate wavelength ranges, extending out to ≈800 Å. In addition, EUVE carried on board a spectrometer for pointed observations following the survey phase of the mission. The spectrometer will be discussed in detail in chapter 6, but half the effective area of its telescope was utilised for a deep survey imager, giving exposure times significantly larger than either the WFC or EUVE all-sky surveys but over a restricted region of sky. The following sections include a detailed description of the payload components drawn from several papers in the scientific literature (see e.g. Bowyer and Malina 1991a,b).

A view of local interstellar space

It is generally accepted that the solar system resides inside a relatively dense local interstellar cloud a few parsecs across with a mean neutral hydrogen density of ≈0.1 cm−3 (e.g. Frisch 1994; Gry et al. 1995). This cloud, the so-called local interstellar cloud (LIC) or surrounding interstellar cloud (SIC), lies inside a region of much lower density, often referred to as the local bubble (figure 7.1). The general picture built up is one where this bubble has been created by the shock wave from a past supernova explosion, which would also have ionised the local cloud. The current ionisation state of the local cloud is then expected to depend on the recombination history of the ionised material, i.e. the length of time since the shock wave passed through. However, if the flux of ionising photons from hot stellar sources is significant, the net recombination rate may be reduced (Cheng and Bruhweiler 1990; Lyuand Bruhweiler 1996). The photometric data from the ROSAT WFC survey have been used to map out the general dimensions of the cavity by Warwick et al. (1993) and Diamond et al. (1995), as already discussed in section 3.9.2. However, this relatively crude interpretation of the observations probably hides greater complexity. For example, several studies of the lines-of-sight towards β and ∊ CMa had already demonstrated the existence of a low density tunnel some 200–300 pc in extent, even before any EUV observations were carried out (e.g. Welsh 1991).

The advent of spaceflight has ushered in a new era of Solar System exploration. Man has walked on the Moon and returned with soil samples. Instrumented probes have descended through the atmospheres of Venus and Mars. The Mariner, Pioneer, Venera, Viking, and Voyager space flight programs have provided opportunities to study the planets from Mercury to Neptune and most of the satellites. Remote sensing investigations have been conducted with unprecedented spatial and spectral resolutions, permitting detailed examinations of atmospheres and surfaces. Even for the Earth, space-borne observations, obtained with global coverage and high spatial, spectral, and temporal resolutions, have revolutionized weather forcasting, climate research, and the exploration of natural resources.

The collective study of the various atmospheres and surfaces in the Solar System constitutes the field of comparative planetology. Wide ranges in surface gravity, solar flux, internal heat, obliquity, rotation rate, mass, and composition provide a broad spectrum of boundary conditions for atmospheric systems. Analyses of data within this context lead to an understanding of physical processes applicable to all planets. Once the general physical principles are identified, the evolution of planetary systems can be explored.

Some of the data needed to address the broader questions have already been collected. Infrared spectra, images, and many other types of data are available in varying amounts for Mercury, Venus, Earth, Mars, Jupiter, Saturn, Uranus, Neptune, and many of their satellites. It is now appropriate to review and assess the techniques used in obtaining the existing information.

The main problem in astrophysics is that of inferring the physical properties of the medium from the observables: the Stokes spectrum. Unfortunately, no in situ measurements can be made of the temperatures, densities, velocities, magnetic fields, and other physical quantities to probe the astronomical object, or at least that portion of the astronomical object where photons come from. Astrophysical measurements are of the physical properties of the (polarized) radiation, not of the celestial object itself. From these measurements, and with the help of some known physics, the astronomer is challenged to infer the properties of the medium that light has passed through. Certainly, we speak loosely when we use the same word measurement for both the process of characterizing light and that of interpreting the observed Stokes spectrum in terms of the medium properties: calibration is neither as easy nor as accurate as in laboratory measurements. The only available “meter” is the RTE, which contains the relationship between the observable (the Stokes spectrum) and the unknowns (the medium physical quantities). More specifically, the link between the medium and the observable lies in the coefficients of the RTE, namely, the propagation matrix and the source function vector.

In this chapter we provide a partial survey of infrared measurements of a number of objects in our Solar System. In view of the large quantity of available data, we have to be very selective in choosing examples. As in other parts of the book, we give space-borne spectrometry preference over broad-band radiometry. Results from lower resolution data can always be simulated by smoothing higher resolution spectra with a corresponding instrument function. To provide a more balanced view of the potential of infrared techniques we occasionally include examples of spectrometric, polarimetric, and radiometric results obtained from ground-based telescopes in addition to spectrometric data from spacecraft.

In Section 6.1 we show the effects of finite spectral resolution and other instrument characteristics on the recording of the emerging radiation field. Infrared data from the terrestrial planets, that is Venus, Earth, and Mars, are treated in a comparative way in Section 6.2. Emphasis is given to an understanding of the physical principles that cause the structure in the measured spectra. The spectra of the giant planets are discussed in Section 6.3, again in a comparative manner. Section 6.4 is devoted to Titan; as a satellite with a deep atmosphere it is in a class by itself. The last section in this chapter (6.5) is concerned with astronomical bodies without substantial atmospheres. Mercury, the Moon, and Io are most interesting examples of this class of objects. The numerical treatment of information retrieval is postponed until Chapter 8.

Now that we have formulated the general RTE for a stratified anisotropic medium in LTE, let us particularize our study to the case of an atomic vapor permeated by a magnetic field, B. For convenience, we shall consider the medium to be isotropic in the absence of an “external” magnetic field. It is thus B that establishes the optical anisotropy by introducing a “preferential” direction.

In order to understand the basic concepts, we start again with the simple Lorentz model of the electron as in Chapter 6 (this time introducing the Lorentz force in the dynamical balance). In this way, the so-called “normal” Zeeman effect gets fully explained. The “anomalous” Zeeman effect, however, needs further results from quantum mechanics that will be summarized later. As the reader may already have realized from the historical introduction, this procedure conforms with historical developments. As in many other branches of physics, a chronological treatment helps in comprehension, although it is not strictly necessary.

In this chapter, we shall see that a single (unpolarized) spectral line in the absence of a magnetic field splits into various Zeeman components, each with a distinct state of polarization that may, of course, vary along the profile.

The Lorentz model of the electron

Let us resume our discussion of Section 6.2 on the Lorentz model.

Since the first edition of this book appeared in print, infrared observations have been responsible for a number of significant new results from many objects in the Solar System. Besides highly sophisticated ground-based measurements, instruments on space probes such as Galileo, Mars Global Surveyor, Vega, Giotto, Phobos-2, the Infrared Space Observatory, and others have produced new data leading to interesting conclusions. Even the spectacular impact of comet Shoemaker–Levy 9 yielded unique information on the atmosphere of Jupiter as well as on the structure of comets. More refined analyses of older data sets have also contributed new insight.

Clearly, an identical reprint of the first edition would have been out of date. To bring the book up to the present state of the art it was necessary to incorporate the latest results. Although discussion of the Solar System bodies has been broadened by including Pluto, comets, and asteroids, the basic format and structure of the book has been preserved. The first four chapters, dealing primarily with fundamental aspects, radiative transfer theory, molecular physics, and modeling of atmospheric spectra, have not been affected by new information. Only minor changes have been made to the text, in some cases to correct errors, in others to clarify certain points. The latest results have been added primarily to Chapters 5 through 9. Some new instrumental techniques needed to be included. More recent information on atmospheric composition and structure had to be compared to older results.

Once one has a solution of the RTE, the most simple procedure of inference can be devised such that a comparison between calculated and observed Stokes spectra suggests modifications in prescribed models of the medium. Iteratively improved models refine the match between observations and theoretical calculations. When the match is good enough, the last model in the iteration is taken as a model of the medium and its characteristic parameters are the inferred parameters of the medium. This trial-and-error method may be useful when the model medium is very simple and contains just a few free parameters. Note that every change of a given free parameter implies an integration of the RTE which is a process requiring a great deal of computer time. If the number of free parameters is large, the manual trial-and-error method can become impracticable, but even automated trialand-error procedures that modify the various parameters randomly (blindly) may not converge to a physically reasonable final model of the medium. The results may even seem reasonable but be greatly in error.

In Chapter 6 we examined planetary spectra using the knowledge that we gained in previous chapters, especially in Chapter 4, which was devoted to simple atmospheric models of radiative processes. By applying physical reasoning, we could extract a considerable amount of information on the conditions that gave rise to the measured spectra. This intuitive method is very important in analyzing data; however, if one desires more precise information in the form of numerical results, the strictly intuitive approach must be augmented by more sophisticated numerical methods. Such methods are the subject of this chapter.

In Section 8.1 we introduce numerical retrieval methods and apply them to atmospheric parameters in general. Section 8.2 is devoted to the retrieval of atmospheric temperature profiles. A large number of different numerical techniques is now available for this task. The retrieval of information on atmospheric composition is the subject of Section 8.3. Again, a wide range of methods must be considered. Cloud parameters and the properties of suspended particulates can also be deduced from infrared measurements. This topic is treated in Section 8.4. The determination of properties of solid surfaces is discussed in Section 8.5, while processes of finding the albedo and the total thermal emission of the Solar System objects are analyzed in Section 8.6.

Were one asked for a concise description of most astrophysical tasks, one possible answer might be ‘understanding the message of light from heavenly bodies’. Light – or electromagnetic radiation – is the astronomer's main (almost his sole) source of information. The statement that nobody can measure the physical parameters of the solar atmosphere, although at first sight shocking, merely calls attention to the fact that astrophysics is an observational rather than an experimental science. This characteristic is often forgotten. We do not measure solar or stellar temperatures, velocities, magnetic fields, etc., simply because we do not have thermometers, tachometers, magnetometers, etc., that would permit in situ measurements of these parameters. Rather, we are only able to measure light. The astronomical parameters are inferred from these measurements, often with the help of some laboratory physics. Thus, the reliability of such astronomical inferences hinges on the accuracy of measurements of light.

This chapter is aimed at understanding how nature and laboratory devices may change the polarization state of light. The transformations of the Stokes parameters are assumed to be linear, i.e., in terms of addition and multiplication by scalars. This is why we are restricted to linear optical systems. The qualifiers quasi-monochromatic and plane will be omitted from now on under the assumption that we are in fact dealing with this type of electromagnetic wave.

Propagation of light through anisotropic media

Changing the polarization state of light means modifying the coherency matrix elements, which in turn require that different components of the electric field vector are acted on differently by the medium. If Ex and Ey suffer the same alteration, a scaling of C is effected, so that the polarization state is unchanged. As a matter of fact, we have seen in the previous chapter how both the linear analyzer and the linear retarder act differently on given components of E. The wave equation (2.1), however, predicts no different behavior for the orthogonal components.