In the first part of the twentieth century, the study of the cosmos (hence cosmology), meaning the study of the universe, was mainly theory. Much of the observational material was beyond the grasp of the telescopes of the time, photography was the only detection medium, and spectroscopy was very primitive. Astrophysics had to thrive by observational studies of the sun, the major planets, and bright stars. Nebulae and clusters had been cataloged by Messier, William Herschel, Caroline Herschel, and John Herschel by the thousands, but few astronomers of the time ever thought they would have the technology to truly study such things in detail.

Lack of observational data has historically been the achilles heel of cosmology. As recently as 1990 Kolb and Turner complained that astrophysics in general and cosmology in particular needed more observations. The present situation is in marked contrast to this view, such that there are now so much data publicly available that it strains our computational ability to analyze it. So many people are currently involved with data reduction and analysis that in the last decadal survey there was a call for the formation of a new area of astronomy and astrophysics called astroinformatics (Borne, 2009).

We start this chapter with several topics that are active parts of this frontier subfield of astronomy and astrophysics, as much of the data in the large databases are galactic and extragalactic. As we shall see, our examination will be limited not by the available data, but by our ability to analyze these data on a personal computer. We begin with an examination of the Hubble law, including evidence of universal acceleration. We then examine galaxies, including radio galaxies and quasars. Finally we look at cosmic structure and evolution.

In this chapter we explore several aspects of dynamics within the Milky Way galaxy, including spiral density waves, the effects of dark matter, and the region near the central black hole. We also examine the effects of galactic rotation on stellar clusters by returning to the Orion Trapezium cluster studied in Chapter 6. Finally we look at the role of gas and dust within the galaxy on astrochemistry.

The existence of dark matter

Because of dust obscuration along the galactic equator, radial velocity studies of bright O and B stars in the 1930s and 1940s were able to discern only three or possibly four spiral arm structures in the part of the galaxy nearest the sun. At the time it was assumed that the derived circular orbit velocities fit Kepler's laws, and from that a “reasonable” mass for the galaxy was obtained. That assumption was accepted without much question. After World War II, with the prediction and subsequent discovery of the atomic hydrogen ground state emission called the 21 cm radiation, a new tool was available for measuring the distribution of matter in the Milky Way and other galaxies.

The proton and electron in an ordinary hydrogen atom are both fermions with a spin of ½. As such they both have magnetic moments that can be either parallel or antiparallel. The energy of the bound electron is slightly higher when parallel rather than antiparallel. This creates an energy difference within the ground state, an effect known as hyperfine splitting. When the electron flips a photon is emitted with a frequency of 1420.4 MHz, which corresponds to a wavelength of 21.1 cm. Unlike visible light, these 21 cm radio waves can penetrate cold (∼10 K) dust clouds, which have predominantly molecular rather than atomic hydrogen, that are found within spiral arms. The atomic hydrogen that produces the 21 cm radiation are found in abundance in the “warmer” (∼100 K) space between the cold molecular clouds. When the Dutch radio astronomer Hendrik C. van de Hulst predicted and subsequently found this radiation, the full extent of the galaxy could be observed and characterized.

It is evident from observation that most of the interstellar medium is permeated with charged particles and permanent magnetic fields. If the ionization in a particular region is complete (no neutral particles) the gas is called a plasma. Thus, interstellar space is always a low density, nearly collisionless, environment where particles may go centuries without encountering a particle of the same kind. It is most certainly dominated by a plasma or at least a highly ionized gas. In this chapter we consider the behavior of light and electric charges in such plasmas in a variety of situations. We will typically use protons and electrons as test particles.

Many treatments of plasmas in astrophysics consider only “cold,” virtually collisionless plasmas, but there are a number of instances, particularly in the vicinity of stars and protostars, where one must examine higher density, higher temperature situations. A cold plasma is one in which the kinetic motion of the protons and electrons generally can be ignored. For warm plasmas, electron and ion temperature becomes a contributing factor and we must take kinetic theory into account via the Maxwell–Boltzmann equation (MBE). Finally we look at two diverse, but actually closely related, applications of plasma theory in astrophysics, the first using pulsars to map the electron density and magnetic field within the Milky Way galaxy, the second being a model of solar wind.

Why another book on astrophysics? Undergraduates and first-year graduate students are deluged with complex information that they are expected to “know” at least qualitatively, and there are plenty of texts that present just such a broad comprehensive survey of astronomy and astrophysics (either observationally or theoretically) and do that just fine. In fact, there are many things being taught in introductory astronomy classes today that 50 years ago appeared only in doctoral theses. But in those same times there were those (Chandrasekhar, Einstein, Hubble, and Spitzer, for example) who, having no access to the powerful computers of today, developed the elegant analytical and observational theories upon which our modern ideas are based. Not only did these early people work with incomplete data sets and poorly understood physical concepts, but they also had to invent their own mathematical and computational methodologies to make their concepts quantitative. The scientific progress of those times was largely the product of sheer intellect from beginning to end.

Astronomy and astrophysics now make such immense strides almost continuously that undergraduate and beginning graduate students are rarely aware of the extraordinary quantitative foundations given to these disciplines during the nineteenth and twentieth centuries, and this knowledge gap widens with each passing day. There are two factors at work here. First, the analytical mathematics used by these early masters was quite above that usually considered suitable for undergraduate instruction, and second the actual computational programming required to produce realistic modern models these days is considered too sophisticated to be meaningfully approached by all but the most advanced undergraduates and beginning graduate students. Yet NASA and other space agencies have often honored the intellectual giants of this era by naming spacecraft after them without explaining to modern students those lines of mathematical/quantitative reasoning that made possible the revolutions in thought of those scientists.

Astrophysics draws upon a wide range of topics in astronomy and physics. Topics as widely ranging as observational techniques, thermodynamics, and general relativity are all central to the material covered in this text. Many readers, particularly undergraduate students, will have only passing experience with some of these foundational concepts. It is impossible for us to provide a comprehensive review of these within the scope of this text, but we begin with a basic overview of the most central topics necessary to approach the astrophysical subjects to be covered. In later chapters we build on these various terms in greater detail. Although we do assume more than a general background in physics and astronomy, we do not provide a comprehensive discussion of the basics, as these can be obtained elsewhere. Instead we present only those more advanced background concepts as needed to the task at hand.

Fundamental stellar properties

One of the central goals of astronomy is the specification of the properties of the sun, stars, and other self-luminous bodies in the universe. Learning about the ranges of these properties and how the quantities characterizing them are determined are a major part of any astronomy course, so we only briefly discuss them here.

Cosmic distance scales

The most fundamental property of a celestial object is its distance from another body. Usually the reference is the sun, the star nearest to Earth. Within the solar system, distances are specified using astronomical units (AU), equal to the mean distance between Earth and sun. Outside the solar system, distances are so large that one immediately switches to another unit, the parsec. The parsec was devised when distances to stars were first measured by the method of trigonometric parallax, and distance was determined from the annual shift of an object’s angular position in the sky.

In Chapters 2 and 3 we looked at atmospheres and interiors of main sequence stars. For such stars we may assume both local thermodynamic equilibrium and the physics of an ideal gas. In this chapter we move beyond these simple assumptions. We begin by looking at atmospheric models in which local thermodynamic equilibrium is not valid, such as the case of expanding atmospheres. We then examine extreme stars such as white dwarfs and neutron stars, in which electron degeneracy plays a central role. As we shall see, this will take us to the limits of classical physics.

Atmospheres beyond local thermodynamic equilibrium

Brief survey of non-LTE situations

In Chapter 2 we developed an atmospheric model that could be approximated by assuming local thermodynamic equilibrium (LTE). Such models work well to describe the photospheres of solar-type main sequence stars. These relatively simple models do not work at all for the chromospheric or coronal regions of a stellar atmosphere. For these regions we must look to non-LTE (NLTE) models.

The most comprehensive treatment of radiative transfer in LTE and NLTE situations is that of Mihalas (1978), and we follow the conventions adopted there in our treatment here. Our present use of NLTE concepts is toward atmospheres that have a systematic (and in some cases differential) macroscopic radial flow. NLTE models are also described in Gray (1992), which outlines how to deal with “turbulence,” which must be handled by direct atmospheric modeling.