Notes on Chapter 20 of "Universe" THE BIRTH OF STARS This chapter describes the interstellar medium and the earliest phases of stellar evolution, up to the production of a stable main- sequence star. Much of the chapter is descriptive. As you examine the various images of astronomical objects (shown in the text, and also shown as slides in class), ask yourself what you are looking at: think about the scale of the object (in parsecs or AU, and its total mass), and why parts of the picture are bright or dark, or red or blue. Concentrate on the first 4 sections of the chapter (20-1 through 20-4; pp. 492-502*). Sections 20-5 through 20-8 are more technical *[(pp. 454-463 in 6th edition] descriptions of specific objects, and here it is sufficient to know what is meant by the terms "T Tauri star" and "giant molecular cloud". At the end of the chapter, the "Key Ideas" and "Review Questions" are at an appropriate level for our class, but the "Advanced Questions" are indeed more advanced. A. GENERAL COMMENTS ON "STELLAR EVOLUTION" Astronomers use the term "evolution" to describe the changes that each individual star goes through in the course of its lifetime. Short- term changes such as pulsations or the development of surface features like spots or flares are not considered evolution; rather, the term refers to irreversible changes in the structure of the star, brought about by the depletion of a nuclear fuel. The "main sequence" is the diagonal band in the HR diagram, running from hot, bright stars to cool, faint stars. Stars on the main sequence differ from one another basically in having different amounts of mass: the upper main-sequence stars (the hot, bright ones) are more massive than the Sun, and the lower main-sequence stars are less massive. The Sun is itself a main-sequence star, one of rather average mass. The main sequence is also a phase of stellar evolution, indeed a phase that lasts most of the star's lifetime. It is the first of the bright, easily observable phases of evolution, starting with the onset of nuclear fusion as an energy source, and ending only when most of the hydrogen in the interior has been changed into helium and is therefore no longer available as a nuclear fuel. When a star first reaches the main sequence -- i.e. when its center becomes hot enough for hydrogen fusion to start -- its movement in the HR diagram stops. The star then spends most of its entire lifetime as a main-sequence star, hardly moving in the diagram. Its position on the main sequence (and hence its spectral type, its absolute magnitude, etc.) is determined by its mass. All main-sequence stars are doing the same thing -- producing energy from the fusion ("burning") of hydrogen into helium at their centers. In the process, a small percentage (0.7%) of the mass involved in the reaction is lost, and is turned into energy. Although most of a star's life is spent on the main sequence, this stage cannot last forever since eventually the hydrogen at its center will be used up. The changes that then occur take it through a number of bright but short-lived phases (red giants or supergiants, various kinds of variable stars, etc.) which are the subject of chapter 21. Here in chapter 20 we consider the phases leading up to the main sequence stage -- the "pre-main-sequence" phase of stellar evolution. This phase is difficult to observe, requiring specialized techniques, since the objects in question emit most of their radiation in the infrared and many of them cannot be detected at all by their optical radiation. Cameras employing infrared detectors have been developed in the last 20 years or so; much of the information in chapter 20 is thus relatively new. B. THE INTERSTELLAR MEDIUM Our Galaxy is a flat, slowly spinning system of immense proportions, at least 30,000 parsecs (100,000 light years) across and containing some 100 billion stars, one of which happens to be our Sun. Although we know today that the Universe contains billions of such galaxies, prior to the 1920s many astronomers considered our Galaxy to be the entire Universe. The visible components of a galaxy are (a) stars, and (b) the interstellar medium, in roughly equal amounts. It has been clear for a long time that the answer to the question, "Where do stars come from?" must be "The interstellar medium," since there are no alternatives. It has also been clear that gravity must be the force that somehow causes an interstellar cloud to pull itself together to form a star. Before we consider how that comes about, we examine the characteristics of the interstellar medium itself. Our Sun lies close to the flat plane of the Galaxy, and as we look around us, the directions where we see a diffuse band of light (the Milky Way) are those in which we see innumerable stars at all distances in this plane. The interstellar medium is also strongly concentrated toward this plane, and therefore we see signs of interstellar material all along the Milky Way, but not so much in other directions. The interstellar matter is very clumpy, and regions where it is relatively highly concentrated are called "interstellar clouds". When these clouds are sufficiently well defined to look like distinct objects, they are called "gaseous nebulae" (or "diffuse nebulae"), or simply "nebulae", from the Latin word for cloud. Many of them have been given picturesque names, but astonomers usually refer to them by their numbers in either of two catalogues: M numbers from Charles Messier's 18th century catalogue of 110 nebulous ("fuzzy"-looking) objects, or NGC numbers from the much more extensive 19th century New General Catalogue of about 7800 non-stellar objects. Both catalogues include a variety of types of objects -- star clusters and galaxies, as well as nebulae -- all of which look indistinct in small telescopes. All nebulae -- and the interstellar medium in general -- consist of two components, "gas" and "dust". It is important to distinguish between these components because they affect light differently, and hence their presence is recognized in different ways. "Gas" is any material that consists of individual atoms or molecules, moving freely. Individual atoms and molecules have the full complement of discrete orbits that their electrons can be in, and they absorb or emit photons of light when an electron moves from one orbit to another. A gas therefore emits and absorbs in SPECTRAL LINES and is completely transparent at other wavelengths. If an interstellar cloud is excited by a nearby hot star, it gives off light in the form of emission lines. If a distant star is seen through an interstellar cloud, the star's spectrum may show extra absorption lines produced in the cloud. Studies of interstellar lines show that the composition of the interstellar gas is generally similar to that of the Sun. "Dust," on the other hand, consists of small solid particles, also called "grains". Although microscopic in size, these particles actually consist of millions of atoms. And because they are solids, they absorb light continuously (at all wavelengths, not in spectral lines), and they emit a continuous distribution of radiation following the law of thermal (blackbody) radiation for the temperature of the grain. We know there is dust in the interstellar medium primarily because it has a major effect on the light of distant stars, making them both fainter and redder. The latter effect ("interstellar reddening") comes about because longer wavelengths get through dusty regions better than shorter wavelengths. Thus all stars seen through a dust cloud are redder than expected for their temperatures (or spectral types), and the amount of reddening tells us the amount of dust. Finally, dust particles are capable of reflecting starlight (just as the "dust tail" of a comet reflects sunlight), and some nebulae are visible because they reflect the light of nearby stars. A cloud of gas can't do this. The gas of the interstellar medium, like that of the Sun and most stars, is mostly hydrogen. However, cool hydrogen is quite difficult to observe. Interstellar hydrogen comes in three forms: 1) H I (neutral hydrogen). A cool interstellar cloud, in which the hydrogen is neutral, is sometimes called an "H I region". At the low temperatures of the interstellar medium, the hydrogen atoms are all in the ground state and therefore cannot absorb in the Balmer lines that we see in stellar spectra. To detect H I directly, we must either use far-ultraviolet observations from space, or, more commonly, a spectral line in the radio part of the spectrum (to be discussed in chapter 25). 2) H II (ionized hydrogen). The space around a very hot star, which emits a lot of energetic UV photons, is called an "H II region" because the hydrogen in that region must be mostly ionized (consisting of free protons and electrons). Any gas cloud within that space will be an emission nebula, giving off an emission-line spectrum by fluorescence. When a proton and electron recombine, the electron is usually in an upper orbit but then drops, in stages, to the ground state, giving off a number of visible and infrared photons. This makes the nebula bright and easily visible. An example is the Orion Nebula. 3) H2 (molecular hydrogen, with a subscript 2). In a cool, dark cloud, the hydrogen is expected to be mostly in molecular form (two protons bound together by shared electrons). Molecular hydrogen itself is very hard to see, but these molecular clouds are known because they also contain dust grains (causing them to look dark) and other molecules such as CO which have emission lines in the microwave region of the spectrum (wavelengths around 1-10 millimeters). We can summarize the types of nebulae as follows: GASEOUS NEBULAE | _________|____________________ | | Bright Dark | _______|____________ | | Emission Reflection Spectrum: emission lines star's spectrum --- Nearby star: very hot bright None Example: Orion Nebula Pleiades Horsehead Dark nebulae are often seen projected against bright objects -- indeed they are hard to see against a dark sky. Thus we find them in rich Milky Way fields, showing up as regions of apparently low star density, or projected against bright nebulae (e.g. the Horsehead Nebula in Orion). Bok globules are small dark nebulae that seem completely black (because of absorption by the dust they contain); they are considered prime sites for future star formation. C. PROTOSTARS AND THE BEGINNING OF STAR FORMATION When an interstellar cloud pulls itself together and starts to contract under its own gravity, it is called a protostar. For the most part, a "protostar" is a theoretical object, because they are very hard to see. They are likely to occur in dusty regions, and owing to their very low temperatures, the radiation they give off is mostly in the infrared. A number of objects detected in recent infrared images are said to be proto- stars or protostar candidates, but it is often hard to be sure if the appellation is appropriate. The so-called "giant molecular clouds" (section 20-7) are large volumes of cool, relatively dense interstellar material which are likely sites for future star formation. They get their name from the way they are discovered. Since molecules have many kinds of energy (electron orbits, vibration and rotation of the atomic nuclei, bending, twisting, etc.), they have thousands of closely-spaced energy levels. Consequently, changes between one energy state and another often result in the emission of low-energy photons which tend to be concentrated in the microwave part of the spectrum, between the far-infrared and radio regions. Microwave technology is a relatively recent development, but now there are several radio telescopes equipped with microwave detectors that can search for spectral lines of molecules. Indeed, thousands of spectral lines, from at least 50 different molecules, including some with as many of 12 or 13 atoms, have now been identified. Most of these molecules have so far been found in only a handful of sources; on the other hand, common molecules like CO have been mapped throughout the Galaxy. What determines when and if an interstellar cloud will contract to form a star? If gravity were the only force acting, all interstellar clouds would have contracted long ago. There must be opposing forces. Some of them are the effects of rotation, and the effects of magnetic fields which run through the interstellar medium. Another is temperature: when a cloud contracts, it gets warmer, and the more energetic collisions between particles at the higher temperatures act to inhibit further contraction. But the cloud radiates away its heat, eventually allowing it to contract further. Thus there is an interplay among several forces which are nearly in balance. Which clouds will contract? Gravity tends to get the upper hand in places where the density of matter is greatest. All interstellar clouds have extremely low densities, but some places such as the Bok globules are denser than others and seem to be good places for the process to start. A number of things can help to trigger the contraction of a cloud. As the Galaxy rotates and the interstellar clouds slosh around, small regions of high density may temporarily form, and then gravity can act to hold those regions together. More interesting is the effect of a supernova explosion -- the violent collapse of a massive star near the end of its life -- which sends shock waves through the interstellar medium, compressing it and possibly touching off a wave of star formation. As a protostar contracts, it releases gravitational energy. In fact, gravitational contraction is the primary source of the energy produced by a star all the way up to the main-sequence phase. This energy warms the star, but it is still cool enough that its radiation is almost entirely in the infrared. The star is now quite bright in its total radiation (its luminosity L), but still faint in absolute visual magnitude. Gradually the star gets smaller and warmer, and it begins to give off some visible light. It keeps getting smaller, hotter, brighter. This progression goes on and on, until finally the star's center reaches a temperature of 10-20 million degrees. Then finally something new happens -- hydrogen at the center starts to fuse (primarily by the proton- proton cycle, in stars like the Sun). Now there is a new source of energy, which heats the interior further, and the resulting increase in thermal pressure causes the contraction to stop. The star is now in a stable state of balance, called hydrostatic equilibrium, in which gravity is balanced by the outward-acting thermal pressure. The star is now a main- sequence star. For a very long time now, it will remain in this state, producing energy at a constant rate and staying at one place in the HR diagram. The duration of this contraction phase depends on the mass: massive stars do everything faster than small stars. A star like the Sun takes about 10 million years to contract, while a massive star might contract in 100,000 years. In any case, the time for contraction is a small fraction of the star's lifetime -- much shorter than the time it will spend on the main sequence. Therefore when we speak of the "age" of a star, we usually mean the time since it reached the main sequence and began hydrogen fusion. D. OBSERVATIONS OF VERY YOUNG CLUSTERS There is evidence that most stars form in bunches, rather than individually. There are regions with lots of Bok globules and small sources that are bright in the infrared -- these seem to be whole clusters of stars in the making. We find star clusters of all ages, and it is likely that most individual stars in the general field were formed in clusters that have dissipated. Small clusters (say, 10-100 stars) do not have enough self-gravity to hold themselves together very long, and they would survive for no more than a few rotations of the Galaxy. Large clusters, on the other hand, can survive essentially forever, and the large clusters that we see are indeed very old. In chapter 21 we will see that the study of star clusters gave important clues regarding the bright phases of stellar evolution. All the stars of a cluster have roughly the same age, but different clusters have different ages; thus we can compare populations of different ages. Also, all the stars of a cluster are at the same distance from us, so that we can compare their absolute magnitudes directly. The main tool we will use to study clusters is, again, the HR diagram, since that lets us see what distribution of star types we have in a group of a certain age. Only a few clusters are so young that they can help us to study pre-main-sequence evolution. One of these is NGC 2264, whose HR diagram is shown in figure 20-17. This diagram should be compared to fig. 20-9, which is a theoretical HR diagram showing "evolutionary tracks" for the pre-main-sequence phases. The idea is that this cluster is so young that many of its stars have not yet reached the main sequence. The most massive stars of the cluster do lie on the main sequence, because they have contracted relatively quickly; but the less massive stars lie in a cloud of points above the main sequence because they have not had time to reach it. We know that NGC 2264 is extremely young for a couple of reasons: (1) Its brightest and most massive star is an extremely hot star, of spectral type O. Such stars have very short lifetimes, only 10 million years or less. If the cluster were much older than that, this massive star would no longer be there as an O-type main-sequence star but would instead be a white dwarf, having completed its evolution. Actually this star seems to be just starting to move off the main sequence -- the first star in this cluster to do so. (2) Many of the smaller (fainter, cooler) stars in the cluster are variable stars of the "T Tauri" type, whose brightness varies in an irregular manner as a sign of their disturbed atmospheres; also, their spectra show unusual emission lines as well as the normal absorption lines, and this is also a sign of a disturbed atmosphere. T Tauri stars with these characteristics are found only in regions of star formation. Interpretation of the HR diagram of the young cluster NGC 2264 will make more sense after we have gone through chapter 21 and studied the HR diagrams of other clusters, of various ages. -- Robert F. Wing January 29, 2001