(Inset A) Disks and jets pervade the universe on many scales. Here, they seem to be a natural result of a rotating cloud of gas contracting to form a star. Matter falling onto the embryonic star creates a pair of high-speed jets of gas perpendicular to the star's flattened disk, carrying away heat and angular momentum that might otherwise prevent the birth of the star. This image shows a small region near the Orion Nebula known as HH1/HH2, whose twin jets have blasted outward for several trillion km (nearly half a light-year) before colliding with interstellar matter. (HH stands for Herbig-Haro, after the investigators who first cataloged such objects.) The next three photos show stellar jets ejected from three different very young stars. Reproduced here to scale, these images collectively depict the propagation of a jet through space.
(Inset B) This image of HH30, spanning approximately 250 billion km, or about 0.01pc, shows a thin jet (in red) emanating from a circumstellar disk (at left in grey) encircling a nascent star.
(Inset C) One of HH34's jets is longer, reaching some 600 billion km, yet remains narrow, with a beaded structure.
(Inset D) HH47 is more than a trillion km in length, or nearly 0.1pc. This photo shows one of its jets plowing through interstellar space, creating bow shocks in the process.
Studying this chapter will enable you to:
Discuss the factors that compete against gravity in the process of star formation.
Summarize the sequence of events leading to the formation of a star like our Sun.
Explain how the formation of a star is affected by its initial mass.
Describe some of the observational evidence supporting the modern theory of star formation.
Explain the nature of interstellar shock waves, and discuss their possible role in the formation of stars.
We now move from the interstellar medium--the gas and dust between the stars--to the stars themselves. The next four chapters discuss the formation and evolution of stars. We have already seen that stars must evolve as they consume their fuel supply, and we have extensive observational evidence of stars at many different evolutionary stages. With the help of these observations, scientists have developed a good understanding of stellar evolution--the complex changes experienced by stars as they form, mature, grow old, and die. We begin by studying the process of star formation, through which interstellar clouds of gas and dust are transformed into the myriad stars we see in the night sky.
Back How do stars form? What factors determine the masses, luminosities, and distribution of stars in our Galaxy? In short, what basic processes are responsible for the appearance of our nighttime sky? Simply stated, star formation begins when part of the interstellar medium--one of the cold, dark clouds discussed in Chapter 18--starts to collapse under its own weight. The cloud fragment heats up as it shrinks, and eventually its center becomes hot enough for nuclear burning to begin. At that point, the contraction stops, and a star is born. But what determines which interstellar clouds collapse? For that matter, since all clouds exert a gravitational pull, why didn't they all collapse long ago? To begin to answer these questions, let us consider a small portion of a large cloud of interstellar gas. Concentrate first on just a few atoms, as shown in Figure 19.1.
Figure 19.1 Motions of a few atoms within an interstellar cloud are influenced by gravity so slightly that their paths are hardly changed (a) before, (b) during, and (c) after an accidental, random encounter.
Each atom has some random motion because of the cloud's heat, even if the cloud's temperature is very low. Each atom is also influenced by the gravitational attraction of all its neighbors. The gravitational force is not large, however, because the mass of each atom is so small. Even when a few atoms accidentally cluster for an instant, as shown in Figure 19.1(b), their combined gravity is insufficient to bind them into a lasting, distinct clump of matter. This accidental cluster would disperse as quickly as it formed. The effect of heat--the random motion of the atoms--is much stronger than the effect of gravity.
Now let's concentrate on a larger group of atoms. Imagine, for example, 50, 100, 1000, even a million atoms, each gravitationally pulling on all the others. The force of gravity is now stronger than before. Would this many atoms exert a combined gravitational attraction strong enough to prevent the clump from dispersing again? The answer--at least under the conditions found in interstellar space--is still no. The gravitational attraction of this mass of atoms is still far too weak to overcome the effect of heat.
We have already seen numerous instances of the competition between heat and gravity. Recall that the temperature of a gas is simply a measure of the average speed of the atoms or molecules in it. The higher the temperature, the greater the average speed, and thus the higher the pressure of the gas. This is the main reason that the Sun and other stars don't collapse. The outward pressure of their heated gases exactly balances gravity's inward pull.
Heat is not the only factor that tends to oppose gravitational contraction. Rotation--that is, spin--can also compete with gravity's inward pull. As we saw in Chapter 15, a contracting cloud having even a small spin tends to develop a bulge around its midsection. As the cloud contracts, it must spin faster (to conserve its angular momentum), and the bulge grows--material on the edge tends to fly off into space. Figure 19.2 illustrates this important feature of rotation. (Consider as an analogy mud flung from a rapidly rotating bicycle wheel.) Eventually, the cloud forms a flattened, rotating disk.
Figure 19.2 A rapidly rotating gas cloud tends to resist contraction. The spin tends to fling matter from the cloud, like mud spinning off the rim of a bicycle wheel. Spin thus competes with the inward pull of gravity. An extended disk of matter forms around the edge of the cloud.
For material to remain part of the cloud and not be spun off into space, a force must be applied--in this case, the force of gravity. The more rapid the rotation, the greater the tendency for the gas to escape, and the greater the gravitational force needed to retain it. It is in this sense that we can regard rotation as opposing the inward pull of gravity. Should the rotation of a contracting gas cloud overpower gravity, the cloud would simply disperse. Thus, rapidly rotating interstellar clouds need more mass for contraction into stars than clouds having no rotation at all.
Magnetism can also hinder a cloud's contraction. Just as Earth, most of the other planets, and the Sun all have some magnetism, magnetic fields permeate most interstellar clouds. As a cloud contracts, it heats up, and atomic encounters become violent enough to ionize the gas. As we noted in Chapter 7 when discussing Earth's Van Allen belts and in Chapter 16 when discussing activity on the Sun, magnetic fields can exert electromagnetic control over charged particles. In effect, the particles tend to become "tied" to the magnetic field--free to move along the field lines but inhibited from moving perpendicular to them. As a result, interstellar clouds may contract in distorted ways. Because the charged particles and the magnetic field are linked, the field itself follows the contraction of a cloud, as indicated in Figure 19.3. The charged particles literally pull the magnetic field toward the cloud's center in the direction perpendicular to the field lines. As the field lines are compressed, the magnetic field strength increases. In this way, the strength of magnetism in a cloud can become much larger than that normally permeating general interstellar space. The primitive solar nebula may have contained a strong magnetic field created in just this way.
Figure 19.3 Magnetism can hinder the contraction of a gas cloud, especially in directions perpendicular to the magnetic field (solid lines). Frames (a), (b), and (c) trace the evolution of a slowly contracting interstellar cloud having some magnetism.
Theory suggests that even small quantities of rotation or magnetism can compete quite effectively with gravity and can greatly alter the evolution of a typical gas cloud. Unfortunately, the interplay of these factors is not very well understood--both can lead to very complex behavior as a cloud contracts, and the combination of the two is extremely difficult to study theoretically. In this chapter we will gain an appreciation for the broad outlines of the star-formation process by neglecting these two complicating factors. Bear in mind, however, that both are probably important in determining the details.
We now return to our original question: How many atoms need to be accumulated for their collective pull of gravity to prevent them from dispersing back into interstellar space? The answer, even for a typical cool (100 K) cloud having no rotation or magnetism, is a truly huge number. Nearly 1057 atoms are required--much larger than the 1025 grains of sand on all the beaches of the world, even larger than the 1051 elementary particles that constitute all the atomic nuclei in the entire Earth. There is simply nothing on Earth comparable to a star.
Back We can best study the specific steps of star formation by considering the HertzsprungRussell (HR) diagram studied earlier in Chapter 17. Recall that an HR diagram is a plot of two key stellar properties--surface temperature (increasing to the left) and luminosity (increasing upward). The luminosity scale in Figure 19.4 is expressed in terms of the solar luminosity (L = 4 x 1026 W). Our G2-type Sun is plotted at a temperature of 6000 K and a luminosity of 1 unit. We can also indicate the size of a star represented by any point on the diagram because of the radiusluminositytemperature relationship. As before, the dashed diagonal lines in the HR diagrams mark stellar radius, allowing us to follow the changes in a star's size as it evolves.
Figure 19.4 The HR diagram is a useful way to summarize the observed properties of both stars and prestellar objects. The diagonal lines correspond to stars of the same radius.
As we saw earlier, most stars plotted on the HR diagram fall along the main sequence. For roughly 90 percent of their lifetimes, stars burn rather quietly and their physical properties do not change much. Data points representing such stable, full-fledged stars remain almost stationary on the HR diagram.
Near the beginning and end of its existence, however, a star's properties change drastically and rapidly. The HR diagram is a useful aid in describing these phases of its life. At each phase of a star's evolution, its surface temperature and luminosity can be represented by a point on the HR diagram. The motion of that point about the diagram as the star evolves is known as the star's evolutionary track. It is a graphical representation of a star's life.
Table 19-1 lists seven evolutionary stages that an interstellar cloud goes through prior to becoming a main-sequence star like our Sun. These stages are characterized by varying central temperatures, surface temperatures, central densities, and radii of the prestellar object. They trace its progress from a quiescent interstellar cloud to a genuine star. The numbers given in Table 19-1 and the following discussion are valid only for stars of approximately the same mass as the Sun. In the next section we will relax this restriction and consider the formation of other stars.
The first stage in the star-formation process is an ordinary dense interstellar cloud, like those studied in Chapter 18. These clouds are truly vast, sometime spanning tens of parsecs (10141015 km) across. Typical temperatures are about 10 K throughout, and densities are usually not much more than 109 particles/m3. Stage-1 clouds contain thousands of times the mass of the Sun in the form of cold atomic and molecular gas.
If such a cloud is to be the birthplace of stars, it must become unstable and eventually break up into smaller pieces. The initial collapse occurs when a pocket of gas becomes gravitationally unstable. Perhaps it is squeezed by some external event, such as the pressure wave produced when a nearby O- or B-type star forms and ionizes its surroundings. Or possibly its supporting magnetic field leaks away as charged particles slowly drift across the confining field lines. Whatever the specific cause, theory suggests that once the collapse begins, fragmentation into smaller and smaller clumps of matter naturally follows, as gravitational instabilities continue to operate in the gas. As illustrated in Figure 19.5, a typical cloud can break up into tens, hundreds, even thousands, of fragments, each imitating the shrinking behavior of the parent cloud and contracting ever faster. The whole process, from a single quiescent cloud to many collapsing fragments, takes a few million years.
Figure 19.5 As an interstellar cloud collapses, gravitational instabilities cause it to fragment into smaller pieces. The pieces themselves continue to collapse and fragment, eventually to form many tens or hundreds of separate stars.
In this way, depending on the precise conditions under which fragmentation takes place, an interstellar cloud can produce either a few dozen stars, each much larger than our Sun, or a whole cluster of hundreds of stars, each comparable to or smaller than our Sun. There is little evidence for stars born in isolation, one star from one cloud. Most stars--perhaps even all stars--appear to originate as members of multiple systems or star clusters. The Sun, which is now found alone and isolated in space, probably escaped from the multiple-star system where it formed after an encounter with another star or some much larger galactic object (such as a molecular cloud).
The second stage in our evolutionary scenario represents the physical conditions in just one of the many fragments that develop in a typical interstellar cloud. A fragment destined to form a star like the Sun contains between 1 and 2 solar masses of material at this stage. Estimated to span a few hundredths of a parsec across, this fuzzy, gaseous blob is still about 100 times the size of our solar system. Its central density is now some 1012 particles/m3.
Even though it has shrunk substantially in size, the fragment's average temperature is not much different from that of the original cloud. The reason is that the gas constantly radiates large amounts of energy into space. The material of the fragment is so thin that photons produced within it easily escape without being reabsorbed by the cloud, so virtually all the energy released in the collapse is radiated away and does not cause any significant increase in temperature. Only at the center, where the radiation must traverse the greatest amount of material in order to escape, is there any appreciable temperature increase. The gas there might be as warm as 100 K by this stage. For the most part, however, the fragment stays cold as it shrinks.
The process of continued fragmentation is eventually stopped by the increasing density within the shrinking cloud. As stage-2 fragments continue to contract, they eventually become so dense that radiation cannot get out easily. The trapped radiation causes the temperature to rise, the pressure to increase, and the fragmentation to cease.
Several tens of thousands of years after it first began contracting, a typical stage-2 fragment has shrunk by the start of stage 3 to a gaseous sphere with a diameter roughly the size of our solar system (still 10,000 times the size of our Sun). The inner regions of the fragment have just become opaque to their own radiation and so have started to heat up, as noted in Table 19-1. The central temperature has reached about 10,000 K--hotter than the hottest steel furnace on Earth. However, the temperature at the fragment's periphery has not increased much. It is still able to radiate its energy into space and so remains cool. The density increases much faster in the core of the fragment than at its periphery, so the outer portions of the cloud are both cooler and thinner than the interior. The central density by this time is approximately 1018 particles/m3 (still only 10-9 kg/m3 or so).
For the first time, our fragment is beginning to resemble a star. The dense, opaque region at the center is called a protostar--an embryonic object perched at the dawn of star birth. Its mass grows as more and more material rains down on it from outside, although its radius continues to shrink because its pressure is still unable to overcome the relentless pull of gravity. After stage 3, we can distinguish a "surface" on the protostar--its photosphere. Inside the photosphere, the protostellar material is opaque to the radiation it emits. From here on, the surface temperatures listed in Table 19-1 refer to the photosphere and not to the "periphery" of the collapsing fragment, whose temperature remains low.
* Note that this is the same definition of "surface" that we used for the Sun in Chapter 16.
As the protostar evolves, it shrinks, its density grows, and its temperature rises, both in the core and at the photosphere. Some 100,000 years after the fragment began to form, it reaches stage 4, where its center seethes at about 1,000,000 K. The electrons and protons ripped from atoms whiz around at hundreds of kilometers per second, yet the temperature is still short of the 107 K needed to ignite the protonproton nuclear reactions that fuse hydrogen into helium. Still much larger than the Sun, our gassy heap is now about the size of Mercury's orbit. Heated by the material falling on it from above, its surface temperature has risen to a few thousand kelvins.
By the time stage 4 is reached, our protostar's physical properties can be plotted on the HR diagram, as shown in Figure 19.6. Knowing the protostar's radius and surface temperature, we can calculate its luminosity. Surprisingly, it turns out to be several thousand times the luminosity of the Sun. Even though the protostar has a surface temperature only about half that of the Sun, it is hundreds of times larger, making its total luminosity very large indeed--in fact, much greater than the luminosity of most main-sequence stars. Because nuclear reactions have not yet begun, this luminosity is due entirely to the release of gravitational energy as the protostar continues to shrink in size and nebular material rains down on its surface.
Figure 19.6 Diagram of the approximate evolutionary track followed by an interstellar cloud fragment prior to reaching the end of the KelvinHelmholtz contraction phase as a stage-4 protostar. (The circled numbers on this and subsequent plots refer to the prestellar evolutionary stages listed in Table 19-1 and described in the text.)
Figure 19.6 depicts the approximate path followed by our interstellar cloud fragment since it became a protostar at stage 3 (which itself lies off the right-hand edge of the figure). This early evolutionary track is known as the KelvinHelmholtz contraction phase, after two European physicists (Lord Kelvin and Hermann von Helmholtz) who first studied the subject. Figure 19.7 is an artist's sketch of an interstellar gas cloud proceeding along the evolutionary path outlined so far.
Figure 19.7 Artist's conception of the changes in an interstellar cloud during the early evolutionary stages outlined in Table 19-1. Shown are a stage-1 interstellar cloud; a stage-2 fragment; a smaller, hotter stage-3 fragment; and a stage-4/stage-5 protostar. (Not drawn to scale.) The duration of each stage, in years, is also indicated.
Our protostar is still not in equilibrium. Even though its temperature is now so high that outward-directed pressure has become a powerful countervailing influence against gravity's continued inward pull, the balance is not yet perfect. The protostar's internal heat gradually diffuses out from the hot center to the cooler surface, where it is radiated away into space. As a result, the overall contraction slows, but it does not stop completely. From our perspective on Earth, this is quite fortunate: If the heated gas were somehow able to counteract gravity completely before the star reached the temperature and density needed to start nuclear burning in its core, the protostar would simply radiate away its heat and never become a true star. The night sky would be abundant in faint protostars, but completely lacking in the genuine article. Of course, there would be no Sun either, so it is unlikely that we, or any other intelligent life form, would exist to appreciate these astronomical subtleties.
After stage 4, the protostar on the HR diagram moves down (toward lower luminosity) and slightly to the left (toward higher temperature), as shown in Figure 19.8. Its surface temperature remains almost constant, and it becomes less luminous as it shrinks. This portion of our protostar's evolutionary path running from point 4 to point 6 in Figure 19.8 is called the Hayashi track, after C. Hayashi, a twentieth-century Japanese researcher. Protostars on the Hayashi track often exhibit violent surface activity. As a consequence, they can have extremely strong protostellar winds, much denser than that of our own Sun. The T Tauri stars discussed in Interlude 191 may well be direct observational evidence of this phase of stellar evolution.
19.8 The changes in a protostar's observed properties are shown by the path of decreasing luminosity, from stage 4 to stage 6, often called the Hayashi track. At stage 7, the newborn star has arrived.
By stage 5 on the Hayashi track, the protostar has shrunk to about 10 times the size of the Sun, its surface temperature is about 4000 K, and its luminosity has fallen to about 10 L. At this point, the central temperature has reached about 5,000,000 K. The gas is completely ionized by now, but the protons still do not have enough thermal energy to overwhelm their mutual electromagnetic repulsion and enter the realm of the nuclear binding force. The core is still too cool for nuclear burning to begin.
Events in a protostar's development proceed more slowly as the protostar approaches the main sequence. The initial contraction and fragmentation of the interstellar cloud occurred quite rapidly, but as the protostar nears the status of a full-fledged star, its evolution slows. The cause of this slowdown is heat--even gravity must struggle to compress a hot object. The contraction is governed largely by the rate at which the protostar's internal energy can be radiated away into space. The greater this radiation of internal energy--that is, the more energy that moves through the star to escape from its surface, the faster the contraction occurs. As the luminosity decreases, so too does the contraction rate.
Some 10 million years after its first appearance, the protostar finally becomes a true star. By the bottom of the Hayashi track, at stage 6, when our roughly 1-solar-mass object has shrunk to a radius of about 1,000,000 km, the contraction has raised the central temperature to 10,000,000 K, enough to ignite nuclear burning. Protons begin fusing into helium nuclei in the core, and a star is born. As shown in Figure 19.8, the star's surface temperature at this point is about 4500 K, still a little cooler than the Sun. Even though the newly formed star is slightly larger in radius than our Sun, its lower temperature means that its luminosity is somewhat less than (actually, about two-thirds) the solar value.
Over the next 30 million years or so, the stage-6 star contracts a little more. In making this slight adjustment, the central density rises to about 1032 particles/m3 (more conveniently expressed as 105 kg/m3), the central temperature increases to 15,000,000 K, and the surface temperature reaches 6000 K. By stage 7, the star finally reaches the main sequence just about where our Sun now resides. Pressure and gravity are finally balanced, and the rate at which nuclear energy is generated in the core exactly matches the rate at which energy is radiated from the surface. The evolutionary events just described occur over the course of some 4050 million years. Although this is a long time by human standards, it is still less than 1 percent of the Sun's lifetime on the main sequence. Once an object begins fusing hydrogen and establishes a "gravity-in/pressure-out" equilibrium, it burns steadily for a very long time. The star's location on the HR diagram will remain virtually unchanged for the next 10 billion years.
Evolution of a 1-Solar-Mass Star
Back The numerical values and the evolutionary track just described are valid only for the case of a 1-solar-mass star. The temperatures, densities, and radii of prestellar objects of other masses exhibit similar trends, but the numbers and the tracks differ, in some cases quite considerably. Not surprisingly, the most massive fragments within interstellar clouds tend to produce the most massive protostars and eventually the most massive stars. Similarly, low-mass fragments give rise to low-mass stars.
Figure 19.9 compares the evolutionary tracks taken by two prestellar objects--one of 0.3 and one of 3 solar masses--with the premain-sequence track followed by the Sun. Notice how all objects traverse the HR diagram in the same general manner, but their luminosities and temperatures differ greatly. Cloud fragments that eventually form more massive stars approach the main sequence along a higher track on the diagram; those destined to form less massive stars take a lower track. More massive objects generally have higher luminosities and surface temperatures at any given evolutionary stage--KelvinHelmholtz contraction, the Hayashi track, or the main sequence--than their less massive counterparts.
Figure 19.9 Prestellar evolutionary paths for stars more massive and less massive than our Sun.
The time required for an interstellar cloud to become a main-sequence star also depends strongly on its mass. Large cloud fragments heat up to the required 10 million K more rapidly than do less massive ones--the most massive fragments contract into stars in a mere million years, roughly 1/50 the time taken by the Sun. The opposite is the case for prestellar objects having masses less than our Sun. Cloud fragments that evolve into low-mass stars are smaller and cooler. Not only do they take a long time to become protostars, but the protostars also take their time changing into full-fledged stars. A typical M-type star, for example, requires nearly a billion years to form, some 20 times longer than it took the Sun to form.
Whatever the mass, the end point of the prestellar evolutionary track is the main sequence. The main sequence predicted by theoretical models, in which stellar properties finally settle down to stable values and an extended period of steady burning ensues, is usually called the zero-age main sequence (ZAMS). It agrees quite well with main sequences observed in the vicinity of the Sun and in more distant star clusters. It is important to realize that the main sequence is itself not an evolutionary track--stars do not evolve along it. Rather, it is just a "way station" on the HR diagram where stars stop and spend most of their lives--low-mass stars at the bottom, high-mass stars at the top.
If all gas clouds contained precisely the same elements in exactly the same proportions, mass would be the sole determinant of a newborn star's location on the HR diagram, and the zero-age main sequence would be a well-defined line rather than a broad band. However, the composition of a star affects its internal structure (mainly by changing the opacity of its outer layers), and this in turn affects both its temperature and its luminosity on the main sequence. Stars with more heavy elements tend to be cooler and slightly less luminous than stars that have the same mass but contain fewer heavy elements. As a result, differences in composition between stars "blur" the zero-age main sequence into a broad band instead of a narrow line.
Some cloud fragments are too small ever to become stars. The giant planet Jupiter is a good example. Jupiter contracted under the influence of gravity, and the resultant heat is still detectable, but the planet did not have enough mass for gravity to crush its matter to the point of nuclear ignition. It became stabilized by heat and rotation before the central temperature became hot enough to fuse hydrogen. Jupiter never evolved beyond the protostar stage. If Jupiter, or any of the other jovian planets, had continued to accumulate gas from the solar nebula, they might eventually have become stars. But virtually all the matter present during the formative stages of our solar system is gone now, swept away by the solar wind.
Low-mass interstellar gas fragments simply lack the mass needed to initiate nuclear burning. Rather than turning into stars, they will continue to cool, eventually becoming compact, dark "clinkers"--cold fragments of unburned matter--in interstellar space. On the basis of theoretical modeling, astronomers believe that the minimum mass of gas needed to generate core temperatures high enough to begin nuclear fusion is about 0.08 solar masses.
Vast numbers of Jupiter-like objects may well be scattered throughout the universe--fragments frozen in time somewhere along the KelvinHelmholtz contraction phase. Small, faint, and cool (and growing ever colder), they are known collectively as brown dwarfs. Our technology currently has great difficulty in detecting them, be they planets associated with stars or interstellar cloud fragments far from any star. We can telescopically detect stars and spectroscopically infer atoms and molecules, but astronomical objects of intermediate size outside our solar system are very hard to see. Interstellar space could contain many cold, dark Jupiter-sized objects without our knowing it. They might even account for more mass than we observe in the form of stars and interstellar gas combined.
Observations of Brown Dwarfs
Back The evolutionary stages we have just described are derived from numerical experiments performed on high-speed computers. Table 19-1 and the evolutionary paths described in Figures 19.6, 19.8, and 19.9 are mathematical predictions of a multifaceted problem incorporating gravity, heat, rotation, magnetism, nuclear reaction rates, elemental abundances, and other physical conditions specifying the state of contracting interstellar clouds. Computer technology has enabled theorists to construct these models, but their accuracy is only partly known, because it is difficult to test them observationally.
How then can we verify the theoretical predictions just outlined? Even the total lifetime of our entire civilization is much shorter than the time needed for a cloud to contract and form a star. We can never observe individual objects proceed through the full panorama of star birth. We can, however, observe many different objects--interstellar clouds, protostars, young stars approaching the main sequence--as they appear today at different stages of their evolutionary cycles. Each observation is like part of a jigsaw puzzle. When properly oriented relative to all the others, the pieces can be used to build up a picture of the full life cycle of a star. By observing premain-sequence objects at many sites in our Galaxy, astronomers have directly verified most of the prestellar stages just described.
Let us now consider in more detail some of the observational pieces that make up the modern picture of prestellar evolution.
Prestellar objects at stages 1 and 2 are not yet hot enough to emit much infrared radiation, and certainly no optical radiation arises from their dark, cool interiors. The best way to study the early stages of cloud contraction and fragmentation is to use radio telescopes to detect the radiation emitted or absorbed by one or more interstellar molecules. Only long-wavelength radiation can escape from these clouds to our telescopes on or near the Earth.
Figure 19.10 shows M20, the splendid emission nebula studied in Chapter 18, along with some of its surroundings. The brilliant region of glowing, ionized gas is not our main interest here, however. Instead, the youthful O- and B-type stars that energize the nebula alert us to the general environment where stars are forming. Emission nebulae are signposts of star birth.
Figure 19.10 The beautiful emission nebula M20 (right) and its dark surroundings (left) provide examples of many phases of star formation. The nebula itself glows because of the energy of the hot young stars embedded in it. The surrounding dark regions show evidence of cloud collapse and fragmentation.
The region surrounding M20 contains galactic matter that seems to be contracting. The presence of (optically) invisible gas there was illustrated in Figure 18.20, which showed a contour map of the abundance of the formaldehyde (H2CO) molecule. These and many other kinds of molecules are widespread in the vicinity of M20, especially throughout the dusty regions below and to the right of the nebula. Their radio emission shows that they are especially abundant near the completely opaque dark region below the nebula. Further analysis of the observations suggests that this region of greatest molecular abundance is also contracting and fragmenting, well on its way toward forming a star--or, more likely, a star cluster.
The interstellar clouds in and around M20 thus provide tentative evidence for three distinct phases of star formation, as shown in Figure 19.11. The huge, dark molecular cloud surrounding the visible nebula is the stage-1 cloud. Both its density and its temperature are low, about 108 particles/m3 and 20 K, respectively.
Figure 19.11 The M20 region shows observational evidence for three broad phases in the birth of a star: (1) the parent cloud (stage 1 of Table 19-1), (2) a contracting fragment (between stages 1 and 2), and (3) the emission nebula (M20 itself) resulting from the formation of one or more massive stars (stages 6 and 7).
Greater densities and temperatures typify smaller regions within this huge cloud. The totally obscured regions labeled A and B, where the molecular emission of radio energy is strongest, are such denser, warmer fragments. Here, the total gas density is observed to be at least 109 particles/m3, and the temperature is about 100 K. The Doppler shifts of the radio lines observed in the vicinity of region B indicate that this portion of M20, labeled "collapsing fragment" in Figure 19.11, is contracting. Less than a light year across, this region has a total mass over 1000 times the mass of the Sun--considerably more than the mass of M20 itself. It lies somewhere between stages 1 and 2 of Table 19-1.
The third star-formation phase shown in Figure 19.11 is M20 itself. The glowing region of ionized gas results directly from a massive O-type star that formed there within the past million years or so. Because the central star is already fully formed, this final phase corresponds to stage 6 or 7 of our evolutionary scenario.
Other parts of our Milky Way Galaxy provide sketchy evidence for prestellar objects in stages 3 through 5. The Orion complex, shown in Figure 19.12, is one such region. Lit from within by several O-type stars, the bright Orion Nebula is partly surrounded by a vast molecular cloud that extends well beyond the roughly 5 × 10 parsec region bounded by the photograph in Figure 19.12(b).
Figure 19.12 (a) The constellation Orion, with the region around its famous emission nebula marked by a rectangle. The Orion Nebula is the middle "star" of Orion's sword. The framed region is enlarged in part (b), suggesting how the nebula is partly surrounded by a vast molecular cloud. Various parts of this cloud are probably fragmenting and contracting, with even smaller sites forming protostars. The two frames at right show some of the evidence for those protostars. (c) A false-color radio image of some intensely emitting molecular sites. (d) A real-color visible image of embedded nebular "knots" thought to harbor protostars.
The Orion molecular cloud harbors several smaller sites of intense radiation emitted by molecules deep within the core of the cloud fragment. Their extent, shown in Figures 19.12(c) and 19.12(d), measures about 1010 km, or 1/1000 of a light year, about the diameter of our solar system. The gas density of these smaller regions is about 1015 particles/m3, much denser than the surrounding cloud. Although their temperature cannot be estimated reliably, many researchers regard these regions as objects well on their way to stage 3. We cannot determine if these regions will eventually form stars like the Sun, but it does seem certain that these intensely emitting regions are on the threshold of becoming protostars.
In the hunt for and study of objects at more advanced stages of star formation, radio techniques become less useful because stages 4, 5, and 6 are expected to display increasingly high temperatures. As the black-body curve of thermal emission from warm protostars and young stars shifts toward shorter wavelengths, these objects should be observable largely in the infrared.
A most interesting object within the core of the Orion molecular cloud was detected by infrared astronomers in the 1970s. Known as the KleinmannLow Nebula, after its discoverers, it is a strong infrared emitter with a luminosity of some 103 L, lying behind the visible nebula. This compact source is outlined by contours in Figure 19.13. Most astronomers agree that this warm, dense blob is a genuine protostar, poised near the end of the KelvinHelmholtz contraction phase, probably around stage 4.
Figure 19.13 The KleinmannLow Nebula, shown here as infrared contours superposed on a visible-light image, is thought to be a young protostar lying within the Orion Nebula.
Until the Infrared Astronomy Satellite was launched in the early 1980s, astronomers were aware only of giant stars forming in clouds far away. But IRAS showed that stars are forming much closer to home, and some of these protostars have masses comparable to our Sun's mass. Figure 19.14 shows a premier example of a solar-mass protostar--Barnard 5. Its infrared heat signature is that expected of an object on the Hayashi track, around stage 5.
Figure 19.14 An infrared image of the nearby region containing the source Barnard 5 (indicated by the arrow). On the basis of its temperature and luminosity, Barnard 5 appears to be a protostar on the Hayashi track in the HR diagram.
The energy sources for some infrared objects seem to be luminous hot stars that are hidden from optical view by surrounding dark clouds. Apparently, some of the stars are already so hot that they emit large amounts of ultraviolet radiation, which is mostly absorbed by a "cocoon" of dust surrounding the central star. The absorbed energy is then reemitted by the dust as infrared radiation. These bright infrared sources are known as cocoon nebulae. Two considerations support the idea that the hot stars responsible for the clouds' heating have only recently ignited: (1) These dust cocoons are predicted to disperse quite rapidly once their central stars form, and (2) they are invariably found in the dense cores of molecular clouds. The central stars probably lie near stage 6.
Protostars often exhibit strong winds. Radio and infrared observations of hydrogen and carbon monoxide molecules, again in the Orion cloud, have revealed gas expanding outward at velocities approaching 100 km/s. High-resolution interferometric observations have disclosed expanding knots of water emission within the same star-forming region and have linked the strong winds to the protostars themselves. As mentioned earlier, these winds may be related to the violent surface activity associated with many protostars (see also Interlude 19-1).
Early in a protostar's life, it may still be embedded in an extensive disk of nebular material in which planets are forming, as discussed in Chapter 15. When the protostellar wind begins to blow, it encounters less resistance in the directions perpendicular to the disk than in the plane of the disk. The result is known as a bipolar flow--two "jets" of matter are expelled in the directions of the poles of the protostar, as illustrated in Figure 19.15. As the protostellar wind gradually destroys the disk, blowing it away into space, the jets widen until, with the disk gone, the wind flows away from the star equally in all directions. Figure 19.16 shows a real example of bipolar flow.
Figure 19.15 (a) When a protostellar wind encounters the disk of nebular gas surrounding the protostar, it tends to form a bipolar jet, preferentially leaving the system along the line of least resistance, which is perpendicular to the disk. (b) As the disk is blown away by the wind, the jets fan out, eventually (c) merging into a spherical wind.
Figure 19.16 (a) This false-colored radio image shows two jets emanating from the young star system HH81-82 (whose position is marked with a cross at center). This is the largest stellar jet known, with a length of about 10,000 A.U. (The colors are coded in order of decreasing radio intensity, red, blue, green.) (b) An idealized artist's conception of a young star system, showing two jets flowing perpendicular to the disk of gas and dust rotating around the star. (See also the chapter-opening photos of more stellar jets.)
Back The subject of star formation is actually much more complicated than the previous discussion suggests. Interstellar space is populated with many kinds of clouds, fragments, protostars, stars, and nebulae. They all interact in a complex fashion, and each type of object undoubtedly affects the behavior of the others.
For example, the presence of an emission nebula in or near a molecular cloud probably influences the evolution of the entire region. We can easily imagine expanding waves of matter driven outward by the high temperatures and pressures in the nebula. As the waves crash into the surrounding molecular cloud, interstellar gas tends to pile up and become compressed. Such a shell of gas, rushing rapidly through space, is known as a shock wave. It can push ordinarily thin matter into dense sheets, just as a plow pushes snow.
Many astronomers regard the passage of a shock wave through interstellar matter as the triggering mechanism needed to initiate star formation in a galaxy. Calculations show that when a shock wave encounters an interstellar cloud, it races around the thinner exterior of the cloud more rapidly than it can penetrate its thicker interior. Thus, shock waves do not blast a cloud from only one direction. They effectively squeeze it from many directions, as illustrated in Figure 19.17. Atomic bomb tests have experimentally demonstrated this squeezing--shock waves created in the blast surround buildings, causing them to be blown together (imploded) rather than apart (exploded). After shock waves cause the initial compression of an interstellar cloud, natural gravitational instabilities may divide it into the fragments that eventually form stars. Figure 19.18 suggests how this mechanism might be at work near M20.
Figure 19.17 Shock waves tend to wrap around interstellar clouds, compressing them to greater densities and thus possibly triggering star formation.
Figure 19.18 An artist's conception of a cloud fragment undergoing compression on the southerly edge of M20, as shock waves from the nebula penetrate the surrounding interstellar cloud.
Emission nebulae are by no means the only generators of interstellar shock waves. At least two other sources are available--the deaths of old stars (planetary nebulae and supernovae, to be discussed in Chapters 20 and 21) and the spiral-arm waves that plow through the Milky Way (to be discussed in Chapter 23). Supernovae are the most energetic, and probably the most efficient, way to pile up matter into dense clumps. These events are relatively few and far between, however, so the other mechanisms may be more important in triggering star formation. Although the evidence is somewhat circumstantial, the presence of young (and thus quick-forming) O- and B-type stars in the vicinity of supernova remnants does suggest that the birth of stars is often initiated by the violent, explosive deaths of others.
Wherever O- and B-type stars have recently formed, we can assume that less massive stars are still in the process of forming. It takes longer for the less massive stars to form, and thus we should not expect to see many A-, F-, G-, K-, or M-type stars directly associated with supernova remnants, provided that the star-formation mechanism really was triggered less than 1 million years ago. The neighborhoods around such remnants are probably vast stellar nurseries--the site of many invisible interstellar cloud fragments and protostars as well as the young, massive stars we see.
This scenario of shock-induced star formation is complicated by the fact that O- and B-type stars form quickly, live briefly, and die explosively. These massive stars, themselves perhaps born of a passing shock wave, may in turn create new shock waves, either through the expanding nebular gas produced by their births or by their explosive deaths. These new shock waves can produce "second-generation" stars, which in turn will explode and give rise to still more shock waves, and so on, as depicted in Figure 19.19. Star formation thus resembles a chain reaction. Other, lighter, stars are also formed in the process, of course, but they are largely "along for the ride." It is the O- and B-type stars that drive the star-formation wave through the cloud. Observational evidence lends some support to this chain-reaction model. Groups of stars nearest molecular clouds do indeed appear to be the youngest, whereas those farther away seem to be older.
Figure 19.19 (a) Star birth and (b) shock waves lead to (c) more star births and more shock waves in a continuous cycle of star formation in many areas of our Galaxy. Like a chain reaction, old stars trigger the formation of new stars ever deeper into an interstellar cloud.
We have seen how a portion of an interstellar cloud can become unstable, collapse, and fragment into stars. Let us take a moment to ask what happens next--not to the newborn stars themselves (that will be the subject of the next three chapters), but to the galactic environment in which they have just formed.
The end result of the collapse process is a group of stars, all formed from the same region of the parent cloud, along with a certain amount of unused gas and dust. How many stars form, and of what type? How much gas is left over? What does the collapsed cloud look like once the star-formation process has run its course? At present, although the main stages in the formation of an individual star (stages 37) are becoming clearer, the answers to these more general questions (involving stages 1 and 2) are still sketchy and await a more thorough understanding of the cloud-collapse process.
In general, the more massive the collapsing region, the more stars are likely to form there. On the basis of observed HR diagrams, we also know that low-mass stars are much more common than high-mass ones. However, the precise number of stars of any given mass or spectral type depends in a complex (and poorly understood) way on conditions within the parent cloud. The same is true of the efficiency of star formation--that is, the fraction of the total mass that actually finds its way into stars--which determines the amount of leftover gas. However, if, as is usually the case, one or more O- or B-type stars form, their intense radiation and winds will cause the surrounding gas to disperse, leaving behind a clump of young stars--a star cluster. Figure 19.20 shows an open cluster in which the gas-dispersal process is almost complete.
Figure 19.20 (a) The Carina emission nebula is seen here in true-color visible light. The region is about 2700 pc from Earth and extends across some 30 pc. (b) An X-ray image of the hottest (O-type) stars clustered near the nebula's center. (c) A new Hubble visible-light image of the nebula's innermost core, known as Eta Carinae--a peanut-shaped region of irregularly scattered gas and dust stretching across about 0.5 pc.
Until recently, the existence of star clusters within emission nebulae was largely conjecture. The stars themselves could not be seen optically because they were obscured by dust. Infrared observations have now clearly demonstrated that stars really are found within star-forming regions! Figure 19.21 compares optical with infrared views of the central regions of the Orion Nebula. The optical image in Figure 19.21(a) shows only the Trapezium, the group of four bright stars responsible for ionizing the nebula. However, the false-color infrared image in Figure 19.21(b) also reveals an extensive cluster of stars (shown here as small crosses of many colors) within and behind the visible nebula. The KleinmannLow Nebula, discussed earlier, can be seen here as the roughly circular blue-green region in the right central portion of the infrared image. The green spot within it is thought to be a dust-shrouded B star just beginning to form its own emission nebula around it. This one image shows many separate stages of star formation.
Figure 19.21 The central regions of the Orion Nebula seen (a) in a short-exposure visible-light image (using a filter that is transparent only to certain emission lines of oxygen) and(b) in the infrared (at roughly the same scale). The visible image shows the nebula itself and four bright O stars, known as the Trapezium. The infrared view, seen here in false color, where red is coolest and white is warmest, shows many cool stars within the nebula that are undetectable in visible light. (c) A short-exposure infrared photograph showing more clearly the four bright Trapezium stars and several faint red stars now emerging from the nebular gas.
For every O or B giant, tens or even hundreds of G, K, and M dwarfs may form. Thus, even a modest emission nebula can give rise to a fairly extensive collection of stars. A typical open star cluster in the plane of our Galaxy, like that shown in Figure 19.22, may measure 10 pc across and contain 1000 or more stars. Less massive, but more extended, clusters are usually known as associations. These typically contain no more than 100 stars but may span many tens of parsecs. Associations tend to be rich in very young stars. Those containing many premain-sequence T Tauri stars are known as T associations, whereas those with prominent O and B stars, such as the Trapezium in Orion, are called OB associations. It is quite likely that the main difference between associations and open clusters is simply the efficiency with which stars formed from the parent cloud.
Figure 19.22 The Jewel Box cluster is a relatively young open cluster in the southern sky. Many bright stars appear in this image, but the cluster contains many more low-mass, less luminous stars. Because some red giants appear among the cluster's blue main-sequence stars, we can estimate the age of the cluster to be about 10 million years.
Eventually, star clusters dissolve into individual stars. One reason is that stellar encounters tend to eject the lightest stars from the cluster, just as the gravitational slingshot effect (see Interlude 6-2) can propel spacecraft around the solar system. At the same time, the tidal gravitational field of the Milky Way Galaxy slowly strips outlying stars from the cluster. Occasional distant encounters with giant molecular clouds also tend to remove cluster stars; a near miss may even disrupt the cluster entirely. As a result of all these influences, most open clusters break up in a few hundred million years, although the actual lifetime depends on the cluster's mass. Loosely bound associations may survive for only a few tens of millions of years, whereas some very massive open clusters such as M67, shown in Figure 19.23, are known from their HR diagrams to be almost 5 billion years old. In a sense, only when a star's parent cluster has completely dissolved is the star-formation process really complete. The road from a gas cloud to a single, isolated star like the Sun is long and tortuous indeed!
Figure 19.23 M67, one of the oldest known open clusters, has survived for almost 5 billion years--an unusually long time for a star system near the plane of the Milky Way Galaxy.
Take another look at the nighttime sky. Ponder all that cosmic activity while gazing upward one clear, dark evening. After studying this chapter, you may find your view of the night sky greatly changed. Even the seemingly quiet nighttime darkness is dominated by continual change.