The Luminosity Oscillations Imager (LOI) is a part of the VIRGO instrument aboard the Solar and Heliospheric Observatory (SOHO) launched on 2 December 1995. The main scientific objectives of the instrument were to detect solar g and p modes in intensity. The instrument is very simple. It consists of a telescope making an image of the Sun onto a silicon detector. This detector resolves the solar disk into 12 spatial elements allowing the detection of degrees lower than seven. The guiding is provided by two piezoelectric actuators that keep the Sun centred on the detector to better than 0.1″. The LOI serves here as an example for understanding the logical steps required for building a space instrument. The steps encompasses the initial scientific objectives, the conceptual design, the detailed design, the testing, the operations and the fulfilment of the initial scientific objectives. Each step is described in details for the LOI. The in-flight and ground-based performances, and the scientific achievements of the LOI are mentioned. When the loop is looped, it can be assessed whether a Next Generation LOI could be useful. This short course can serve as a guide when one wishes to propose a space instrument for a new space mission.

Absorption lines in stellar spectra show differences in shape and strength according to the physical conditions in the star's atmosphere. Some of the toughest and most fascinating problems arise in the study of the interplay of the line absorption with the temperature, pressure, radiation, and magnetic and velocity fields of the gas. We are not yet able to calculate the full interlinking of these variables. On the other hand, the panoramic view of spectral-line behavior can be understood in relatively simple terms which we shall presently describe.

The most fundamental point to bear in mind is that the strength of line absorption depends on the number of absorbers producing that absorption. Thus the atomic level populations are of primary concern. But since the number of absorbers also means along the line of sight through the visible depths of the atmosphere, the path length is equally important. Specifically, if the continuous absorption is strong, the path length will be short and vice versa. In effect, the depth of the atmosphere changes with the amount of continuous absorption. In this way, the ratio of the line absorption to the continuous absorption is seen to be the main factor to consider.

One of the goals of stellar atmosphere studies is to understand the various line profiles and line strengths shown by stars. Another is to use our knowledge of line behavior to interpret the fundamental properties of stars, for example, the measurement of effective temperature, surface gravity, radii, and chemical composition. Our attention is directed to the first of these goals in this chapter, while the others are taken up in subsequent chapters.

As a precursor to calculating the transfer of radiation through a model stellar photosphere, we now look at the continuous absorption coefficient. The wavelength dependence of the continuous absorption coefficient shapes the continuous spectrum emitted by the star: more absorption, less light. The strength of spectral lines also depends on the continuous absorption; more continuous absorption means a thinner photosphere with fewer atoms to make spectral lines. Consequently, we must know kν in order to match model-computed spectra to real stellar spectra. However, before we can compute the theoretical spectrum, we need to compute the model on which it is based, and that too can require kν. Specifically, if we invoke radiative equilibrium to find T(T0), then kν is needed from the very beginning. Furthermore, it then becomes particularly important to know kν in those spectral regions carrying most of the flux. In hot stars, this means the ultraviolet region, in cool stars, the infrared region. We escape this step if T(T0) is determined empirically, say by scaling the solar temperature distribution; the main need for kν reverts to the computation of the spectrum. Since scaled temperature distributions are emphasized in this book, kν in the visible window is the main focus of this chapter.

What is a model photosphere, and why build one? It would seem logical to take our stellar observations and deduce from them the physical conditions existing in the atmosphere of the star – somewhat like a parallax measurement yields the distance to a star. Alas, the formation of the stellar spectrum involves many physical variables, and a rigorous deductive interpretation cannot generally be made. Instead we hypothesize a model through which we organize and relate the information conveyed in the starlight. The model resembles a scientific theory in that it is constructed on the basis of our observations and known physical laws. We may then gather additional observations to test our model much like we would test a theory. The model is modified and improved as information is added. When our model closely reproduces all the available observations, we begin to feel that our model is worthy of some trust. Then properties associated with the model, or deduced from further application of the model, are associated with the star, properties such as effective temperature, surface gravity, radius, chemical composition, or rate of rotation.

The model photosphere consists of a table of numbers giving the source function and the pressure as a function of optical depth for an assumed chemical composition. Additional columns may be added to the table depending on the use intended for the model.

Since the 1960s, instruments on Space Physics missions have changed and improved in many respects, but the basic physical parameters that we need to measure have remained the same. The requirement on any Space Physics mission is still to make measurements which are as extensive as possible of all the parameters of space plasmas: the distribution functions of all the constituents in the plasma populations; the DC and AC magnetic and electric fields; and the distribution functions of energetic particles species. All these parameters and distribution functions need to be measured with high spatial, directional and temporal resolution.

These lectures rely on extensive experience building magnetometers, energetic particle detectors, as well as on-board data processors and power supply and power management systems for space instruments. They provide an overview of the kind of instrumentation in which Europe has acquired considerable expertise over the years and which will be continued in future missions.

When I gave this talk in the Canary Islands Winter School of 2003, it was obvious that the interest of the audience was about how to make a successful proposal rather than finding out about the developing phases of a space mission. Unfortunately, I do not know how to make a 100% successful proposal. Success depends on a combination of bright ideas, creativity, timely response to the needs of a large scientific community, adequate system knowledge and, certainly, a bit of good luck. This presentation aims to make young scientists acquainted with the phases and challenges encountered in new space science missions. For that purpose these notes are organized in two sections. The first one establishes the phases of a mission, that is the process of carrying through a generic science project, while the second deals with the actual role of scientists in the whole process. Other talks in the Winter School focused in the science and the experiments that might be done, on how we can increase our knowledge of the Universe by means of space technologies. Here, we try to help making these, as well as other new ideas, real space science experiments.

The steps needed to define a successful space science mission are numerous. The science drivers, the unique advantages this mission provides over past missions or earth-based experiments, and the payload that it includes are the key factors to guarantee its success. Finding the required information on such topics is not so straightforward, especially as they are usually outside the scope of undergraduate courses. The 2003 Canary Islands Winter School of Astrophysics aimed at providing a focused framework that helps fill this need. Space agencies follow a necessarily complex path towards the selection of a specific mission, as required by the enormous costs that are associated with space activities. The steps towards its completion are elaborate and require careful assessment at every stage. The orbit that will be used and the requirements that are imposed have impacts on the science and the mission budget. Thus, knowing how to make the best use of propulsion technologies and gravity helps from solar system bodies plays a crucial role. The first two chapters of this book cover all these topics and illustrate the complexity of defining space missions as well as how and where look for help (i.e. other than the rarely receptive funding agencies).

The instruments on-board will in the end make the science that has driven the mission. How the science questions translate into specific requirements, and then, into the actual instruments are crucial aspects in the definition of the payload.

These lectures cover the principles of remote sensing instrumentation as commonly used for missions to solar system bodies. From the basic physical principles, an introduction to real instruments is provided. Particular attention is paid to the airborne visible infrared imaging spectrometer (AVIRIS) hyperspectral imager, Cassini-visual and infrared mapping spectrometer (VIMS) and the visible and infrared thermal imaging spectrometer (VIRTIS) series. We conclude with a review of the in situ science provided by landers, rovers and other surface elements.

In this chapter, we concern ourselves primarily with the spectroscopic effects of stellar rotation, including the line broadening caused by rotation, techniques for extracting the rotation rate from the broadening, and some of the results. Many of the analysis tools are the same as those discussed in Chapter 17.

Stellar rotation is the driving force for diverse phenomena in stellar atmospheres such as circulation currents, mass loss, and magnetic field generation and its offshoots: starspots, flares, chromospheres and coronae, and activity cycles. The existence of stellar rotation comes as no surprise. Since we believe stars form out of interstellar clouds, and clouds collide with one another and are subject to other torques such as those induced by galactic rotation, we would predict rather large stellar rotation on the basis of the conservation of angular momentum. Moment of inertia scales roughly as the square of the linear dimension. With 44 million solar radii per parsec, cloud rotation should be magnified by something like 15 orders in coming to stellar dimensions. Some stars rotate near the “break-up” velocity, where the gravitational acceleration is comparable to the centripetal force at the equator, but even these rates are small compared to what is expected. Apparently dissipation of angular momentum is an integral part of star formation.

When we look at stellar spectra, we soon recognize the presence of many chemicals. The lines of hydrogen dominate the photospheric spectra of hot stars, but with declining temperature, thousands of lines from other species grow stronger while the lines of hydrogen weaken. One of the tasks of the stellarphotosphere analyst is to disentangle the effects of temperature, pressure, turbulence, and so on from the effects of chemical composition. In broad strokes, the results of chemical analyses tell us that hydrogen is overwhelmingly the most abundant element, comprising ≈90% of the atoms in a normal stellar photosphere, and helium comes next with ≈10%. The remaining elements, often referred to as “metals,” comprise a sprinkling – like salt in a bowl of broth – adding savor and interest. Most of the spectral lines are caused by these low-abundance elements.

Besides our natural curiosity to know what stars are made of, we are motivated to perform chemical composition studies on other fronts. To be more specific, almost any study in stellar atmospheres presupposes some chemical composition. Chemical evidence tells us about the nuclear reactions taking place in stars, internal mixing of material, penetration depths of convection zones, diffusion and gravitational settling, or accretion of material from interstellar space. Lithium and carbon isotope abundances are signatures of the evolutionary stage of a star. The evolution of our galaxy is traceable in part through the differences in chemical composition among stars, since we expect stars formed at different times and in different places to retain information concerning the composition of the material from which they formed.

Observations of stellar photospheres require instruments for collecting and analyzing light. Low-resolution spectrographs are used for measuring the continuous spectrum, while high-resolution spectrographs are needed to measure the spectral lines. In this chapter, we delve into several basic aspects of spectrographs and diffraction gratings, then turn our attention briefly to interferometers, and finally to telescopes. Special features and application of spectrophotometric equipment are discussed separately in Chapter 10 (continua) and Chapter 12 (lines). Chapter 4 is about light detectors.

Spectrographs: some general relations

The basic astronomical spectrograph is shown in Fig. 3.1. It consists of an entrance slit placed at the focus of the telescope, a collimator that intercepts the divergent telescope beam, a dispersing element (prism or grating), and a camera that focuses the dispersed light onto a detector. Since the purpose of the collimator is to make the divergent beam parallel, the distance between the slit and the collimator is the focal length of the collimator, Fcoll. Similarly, the distance between the camera and the focused spectrum is the focal length of the camera, Fcam. The distances between the collimator, disperser, and camera affect the detailed design and optimization of the spectrograph, but do not matter for the basics we are considering at the moment. We concentrate on plane diffraction gratings as the dispersing element since these are almost universally used in astronomical spectrographs.

The remarkable nature of stars is transmitted to us by the light they send. The light escapes from the outer layers of the star – called, by definition, the atmosphere. The complete atmosphere of a star can be viewed comprehensively as a transition from the stellar interior to the interstellar medium. And yet almost the whole visible stellar spectrum comes from a relatively thin part called the photosphere. Obviously we cannot disconnect the photosphere from the adjacent portions of the atmosphere, but in actual fact it is the only region we can study extensively for most stars. It is for this reason that the photosphere has taken its place as the central theme of this book.

Several books have appeared during the last decade dealing with the theory of stellar atmospheres. These works are for the most part excellent. It is to the material largely omitted by these books that the present treatise is directed. My students and I have felt for some time the need of a book that presents the basics of the field through the eyes of an observer and analyzer of stellar atmospheres.

An introduction to a subject, in my opinion, should be presented in a way that can be understood by a reader who has not studied the topic before. It follows that the material should be presented in as simple and straightforward a manner as possible. The Fourier transform (as covered in Chapter 2) is a unifying theme helping to accomplish this aim.

There is hardly a phase of modern astrophysics to which Fourier techniques do not lend some insight or practical advantage. Fourier concepts prove useful in the context of line absorption coefficients, the analysis of line profiles, spectrograph resolution, telescope diffraction, and the study of noise. In these and other applications, convolutions appear in the physics of the situation. Usually it is much easier to visualize a product of functions in place of their convolution and this can be done with Fourier transforms through the convolution theorem.

This chapter forms an introduction to Fourier transforms for those unfamiliar with them and a useful refresher for those who have studied them in past years. The treatment is highly abbreviated, but covers all the concepts used in the remainder of the book. Those wishing to learn the material in a more rigorous and extensive way are referred to the books dealing specifically with Fourier transforms, for instance, the books of Jennison (1961), Bracewell (1965), and Gaskill (1978).

The definition

The Fourier transform of a function is a specification of the amplitudes and phases of sinusoidals which, when added together, reproduce the function. Only one-dimensional functions are treated here. Expansion to two or higher dimensions can be done by the reader with modest effort.

Wonderful growth has occurred in our understanding of stellar photospheres during the 15 years since the appearance of the first edition of “Photospheres.” I have managed to retain the same chapter names and the general plan of the first edition, and many of the equation numbers are also the same. But a significant portion of the material is new or revised. A revolution in light detectors has given us hundreds of times greater efficiency in measuring stellar spectra; Chapter 4 on detectors has been re-done. The astronomical literature is burgeoning with new results on the structure of photospheres, chemical abundances, radius measurements, stellar rotation, and photospheric velocity fields. Many of these results have been incorporated in this second edition, of course. At the same time, I stayed with my original purpose of making this volume an introduction to the subject. Unhappily, this means leaving out numerous exciting topics. My book Lectures (Gray 1988) takes up some of these, and it is recommended as a second installment, after the material in “Photospheres” has been mastered.

More than ever, the reader should keep in mind the fundamental nature of the stellar photosphere: of interest in its own right, with marvelous and intriguing physics, yet the link between the interior and chromospheres, coronae, and interstellar surroundings, and the source of most of our basic information about stars and stellar systems.