The Milky Way and Beyond (28 page)

In H II regions, hot gas also arises from the stellar winds of the exciting stars. These winds create a large cavity or bubble in the denser, cooler gas originally surrounding such a star. In the interior of the bubble, the radially flowing stellar wind passes through a transition in which its radial motion is converted into heat. The hot gas then fills most of the cavity (perhaps 90 percent or more) and serves to separate the filaments of the warm, comparatively dense H II region. Within the condensations of visible plasma, there are neutral globules in which the gas is quite cold (about 100 K) but is dense enough (typically, 10,000 atoms per cubic cm) to have about the same pressure as the hot and warm materials. In short, an H II region is much more complicated than its visual radiation would suggest.

H II regions are almost always accompanied by molecular clouds on their borders. The Orion Nebula, for example, is merely a conspicuous ionized region on the nearby face of a much larger dark cloud; the H II region is almost entirely produced by the ionization provided by a single hot star, one of the four bright central stars (the Trapezium) identified by Dutch astronomer Christiaan Huygens in 1656. The shape of the Orion Nebula appears at visible wavelengths as irregular. However, much of this seeming chaos is spurious, caused by obscuration of dust in dark foreground neutral material rather than by the actual distribution of ionized material. Radio waves can penetrate the dust unhindered, and the radio emission from the ionized gas reveals it to be quite circular in shape and surprisingly symmetrical as seen in projection on the sky. The foreground dark material obscures about half the ionized nebula.

An H II region on the outer edge of a large molecular cloud can induce star formation. For instance, behind the bright Orion Nebula, deeper within the dark cold Orion molecular cloud, new stars are being formed today. At present, none of the new stars is massive and hot enough to produce its own H II region, but presumably some of them eventually will be. When an H II region is produced from cold molecular gas by the formation of a hot star, the temperature is raised from roughly 25 to 8,000 K, and the number of particles per cubic centimetre is almost quadrupled because each H
₂
molecule is split into two ions and two electrons. Gas pressure is proportional to the product of the temperature and number of particles per cubic centimetre (regardless of their
mass, so electrons are as important as the much heavier ions). Thus, the pressure in an H II region is some 800 times the pressure of the cold gas from which it formed. The excess pressure causes a violent expansion of the gas into the dense cloud. Rapid star formation may occur in the compressed region, producing an expanding group of young stars. Such groups, the so-called O Associations (with O stars) or T Associations (with T Tauri stars), have been observed. The component stars simultaneously generate extremely fast outflows from their atmospheres. These winds create regions of hot, tenuous gas surrounding the association. Eventually the massive stars in the association explode as supernovae, which further disturb the surrounding gas.

This picture of the evolution of H II regions and molecular clouds is one of constant turmoil, a few transient O stars serving to keep the material stirred, in constant motion, continually producing new stars and churning clouds of gas and dust. In this way some of the stellar thermonuclear energy is converted into the kinetic energy of interstellar gas. This process begins just after the formation of the massive star that will power the mature H II region. The star begins producing copious amounts of ultraviolet radiation, converting the surrounding H
₂
to atomic hydrogen, and then ionizing it to very dense high-pressure H
⁺
.

The very dense, very young H II regions within molecular clouds are called “ultracompact” because of their small sizes and high densities. These nebulae are observed only at the wavelengths of radio and far-infrared radiation, both of which are able to penetrate the thick dust in the clouds. They are extremely bright at wavelengths of 50 micrometres (.002 in). There are about 2,500 in the Milky Way Galaxy, representing 10 to 20 percent of the total O-type star population. Usually only a light-month in size, 100 times smaller than a typical H II region, they show densities in the ionized region of 10
⁵
hydrogen atoms per cubic cm. They cannot be at rest with respect to the surrounding gas; if they were, the immense pressure exerted by their dense hot gas would cause a violent expansion. (Their lifetimes would be only about 3,000 years—exceedingly short on an astronomical timescale—and not nearly as many could be seen as the number observed by astronomers.) Rather, their gas is kept confined because they are moving through the surrounding cloud at speeds of about 10 km (6 miles) per second, and what is observed is the cloud of freshly ionized gas ahead of them that has not yet had time to expand. The ultracompact H II regions leave behind a trail of ionized material that is not as bright as the confined gas ahead of them. This trail gradually fades as it recombines after the ionizing star has passed. The radio radiation is produced by the ionized gas, but the far-infrared
radiation is emitted by the 5,000 solar masses of surrounding dust warmed by the luminosity of the embedded star.

The most energetic H II regions within nearby galaxies have over 1,000 times more ionizations per second than does the Orion Nebula, too many to be provided by a single star. Indeed, there are clumps of ionized gas ionized by tight groupings of single stars that are embedded in rather diffuse material. These objects are more than 10 times as luminous as any in the Milky Way Galaxy and are about 200 light-years in diameter. If they were located at the Orion Nebula, they would cover the entire constellation of Orion with brightly glowing gas. These supergiant nebulae are more than 10 times as luminous as any in the Galaxy.

The entire Local Group—the cluster of galaxies consisting of the Milky Way Galaxy, the great spiral galaxy in Andromeda, the smaller spiral in Triangulum, and more than 50 other stellar assemblages—contains but one supergiant nebula, the object called 30 Doradus, in the Large Magellanic Cloud. It contains a stellar cluster called R136, the source of most of the energy radiated by the nebula. This grouping consists of dozens of the most massive known stars of the Milky Way Galaxy, all packed into a volume only a thousandth of a typical stellar spacing in size. How such a cluster could form is a fascinating puzzle. There are other supergiant nebulae outside the Local Group, some of which radiate 10 times the energy of 30 Doradus.

The chemical composition of H II regions (the numbers of atoms of each chemical element, relative to hydrogen) can be estimated from nebular spectra. Each element is found in more than one stage of ionization, so the first step is to use the emission-line strengths of each stage of ionization, relative to those of the hydrogen lines, to obtain the abundance of that particular stage of ionization. All abundant elements have some stages of ionization that produce observable emission lines. On the other hand, some elements, such as argon, sulfur, and carbon, have important ions that do not show easily observable lines. Elaborate computer calculations predict the ionization structure of gas ionized by a hot star whose temperature is determined by its spectrum. The calculations then provide predictions of the abundances of the invisible ions, relative to the observable, and the total elemental abundance follows.

The main difficulty with this straightforward procedure is that there are two methods for determining the observed ionic abundances; each should be reliable, but they give quite different answers. By far the easier method of determining ionic abundances is to observe the bright lines produced by collisions between the ion and energetic electrons. The brightest lines from this process in the entire spectrum of H II regions arise from either O+ or O
⁺²
. These bright lines (and others such as N
⁺
) are the basis of the abundance determinations in other galaxies.

Alternatively, the ionic abundances can be determined from the very faint emission lines that follow recombination, the process by which the higher stage of ionization captures an electron (usually at low energies) into a high level of the ion. Following recombination, there is a cascade from the high energy levels to the ground state, with photons in the observed emission line being emitted at each downward transition. These emission lines are fainter than the hydrogen emissions by roughly the ratio of abundances, which is more than 1,000, so only in bright nebulae can this method be
used. However, modern spectrographs on large telescopes have provided strengths of these faint recombination lines for many objects, with good agreement among observers. Checks are provided by comparison of several lines of the same stage of ionization. The relative strengths of the observed lines agree well with the expectations of the cascading process in the excited ions.

The results are interesting and controversial. For H II regions that are bright enough for the faint heavy-element recombination lines to be measured, the carbon, nitrogen, neon, and oxygen abundances from recombination lines are uniformly about 1.8 times those from the collisionally excited lines. A common interpretation is that there are strong temperature fluctuations within the nebulae. In the warmer regions the collisionally excited lines are strongly overproduced per heavy ion, so fewer heavy ions are needed to account for the observed line strengths. The hydrogen lines are hardly affected by the postulated temperature fluctuations. The temperature fluctuations, which must be large (about 20 percent of the average), are unexplained. Turbulence and magnetic fields are prime suspects.

The nebular temperature can be estimated directly from collisional lines alone by comparing emission lines from high-energy levels, populated by collisions, with lines from lower levels. This process can be carried out at various places on the sky within each nebula, and large-scale temperature fluctuations would appear as variations from place to place. However, no such variations have been convincingly detected. Some astronomers have proposed that there are chemical inhomogeneities within nebulae that give rise to the differences between the abundances derived from recombination and collisionally excited lines, but how such variations could be maintained is unexplained.

Nevertheless, abundances are estimated on the basis of a simple interpretation of the mysterious postulated temperature excursions, but no better procedure has been suggested to deal with the startling discrepancies in the derived abundances. The other class of ionized nebulae, the planetary nebulae, show the same effect. The local estimated abundances (say, for the Orion Nebula) are roughly solar. The abundances of heavy elements per million hydrogen atoms are 500 for carbon, 80 for nitrogen, 600 for oxygen, and 100 for neon. There is a gradient of these abundances within the Milky Way Galaxy. At a distance halfway to the centre, 12,000 light-years inward, they are 50 percent larger than locally. Beyond Earth the gradient seems to persist, but there are very few observations. Helium is about 0.1 times as abundant as hydrogen, by number of atoms, throughout the Milky Way Galaxy.

Except for a few cases, compositions of nebulae in galaxies outside the Milky Way Galaxy are measured by collisionally excited lines. The Large Magellanic Cloud has compositions that are uniformly about one-half those of the local Milky Way for oxygen, neon, argon, and
sulfur and are one-quarter for carbon and nitrogen. It appears that the first group of elements must be manufactured together, presumably in massive stars, and ejected together into the interstellar gas. Stars of a different (probably lower) mass must produce carbon and nitrogen. Planetary nebulae suggest the same scenario.

The abundance of helium in nebulae has received considerable attention because the helium content of the oldest objects provides clues to the origin of the universe. The value cited above for the Orion Nebula is in agreement with the predictions of the big-bang model.

Some nebulae, called planetary nebulae, are expanding shells of luminous gas expelled by dying stars. Observed telescopically, they have a relatively round compact appearance rather than the chaotic patchy shapes of other nebulae—hence their name, which was given because of their resemblance to planetary disks when viewed with the instruments of the late 1700s, when the first planetary nebulae were discovered.

There are believed to be about 20,000 objects called planetary nebulae in the Milky Way Galaxy, each representing gas expelled relatively recently from a central star very late in its evolution. Because of the obscuration of dust in the Galaxy, only about 1,800 planetary nebulae have been cataloged. Planetary nebulae are important sources of the gas in the interstellar medium.