No one has ever observed stars go through their life cycle, but astronomers can observe many stars at different stages in their life cycles. In addition, astronomers can calculate what would happen to stars under various conditions, and attempt to match the predictions against actual stars. The overall outlines of stellar evolution are probably accurately known, but there are many unanswered questions, some of them big ones, and very likely there are surprises waiting for astronomers as well.
Stars are believed to form when clouds of interstellar gas and dust start to contract under the influence of gravity. In interstellar space, far from other stars, a cloud of gas can be very thin and still be dense enough to begin contracting. Many astronomers believe violent explosions of older stars create shock waves that help start the contraction process, at the same time enriching the gas in heavy chemical elements. In all likelihood, a cloud will condense into many stars and form a star cluster.
Initially, a cloud destined to become a star will be roughly spherical. The cloud will almost certainly have a slight rotation, simply because of random gas motions in the cloud as it started contracting. As the cloud contracts, its rotation will speed up, causing the cloud to become disk-shaped. The cloud is still far larger than our solar system, and its outer reaches very thin. Most of the mass of the cloud falls to the center, adding energy to the center and heating it up. The center of the cloud becomes a protostar, emitting mostly infrared radiation. Finally the temperatures and pressures within the star reach the point where nuclear reactions begin, and the star "turns on". Matter in the surrounding disk may accrete to form planets, or, if the condensations are massive enough, companion stars.
Many, perhaps all stars, have companions. Many stars are binary stars or multiple stars, with several stars orbiting around one another. Some multiple stars may form when a rapidly rotating protostar becomes unstable and splits; others may form as separate protostars. Our nearest neighbor, Alpha Centauri, is a triple star, with a primary star much like our sun, a smaller and cooler secondary star orbiting about as far away as Uranus orbits our sun, and a very faint companion many times farther away from the primary than pluto is from the sun. Binary stars are of great value to astronomers because the orbital periods of binary stars depend on the masses of the stars, and binary stars enable astronomers to determine the masses of stars
Does the sun have a companion star? There has been speculation from time to time that the sun might have a companion, most recently the "nemesis" hypothesis. The "nemesis" hypothesis argued that an undiscovered companion of the sun periodically caused comets to sweep through the solar system, triggering among other things the extinction of the dinosaurs. If any undiscovered companion exists, it must be very faint, not very massive, and very far away. The odds are very much against such a companion escaping detection this long.
The sun has planets. From what we know of the formation of the solar system, it seems very likely that many if not most solitary stars have planets, formed from the material that did not condense into the central protostar. Possibly widely-separated multiple stars like alpha centauri have planets as well. Whether exotic multiple stars with many components, or very close binary systems can have planets is uncertain. Even if planets form around such stars, close encounters with the stars may fling the planets into a star or out of the star system altogether. A number of stars have disks of solid matter orbiting them that are believed to be similar to the disk from which the planets of our Solar System formed. Such disks have been termed proplyds (Protoplanetary Disks).
Locating other planetary systems is a very great challenge. One approach is to detect very tiny, regular variations in the positions of stars. As a planet orbits a star, the star and planet both orbit around their center of mass; if the star and planet were joined by a beam, the point where the beam would balance is the center of mass. Tiny motions of the star can also be detected using the Doppler Effect. The motions are very tiny: it would be hard to detect the effect of Jupiter on the sun from a nearby star. Another approach is to block the light of the star and search for the reflected light from planets. All of these techniques are extreme challenges to existing technology but are becoming more feasibile as instruments improve. Beginning in 1993 evidence for other planetary systems has rapidly accumulated. Some of these newly-discovered objects are so massive they may be so-called brown dwarfs, on the borderline between large planets and tiny stars.
The newly-discovered objects around other stars are very unusual. None of the probable other solar systems look much like our own. Many of the planets are much more massive than Jupiter and some of these objects orbit surprisingly close to their suns. It appears that our solar system may be unusual.
Are there objects midway between planets like Jupiter and small stars? The fainter stars are, the more numerous, so that we might expect such objects, nicknamed brown dwarfs to be very common. None have been conclusively detected, yet, even though their infrared radiation would be easily detectable. It appears that there is a sharp dividing line between solitary and multiple stars, for reasons still unclear. Some of the newly-discovered objects around other stars are so massive they may be brown dwarfs.
Once stars begin to shine, they assume a position on the main sequence and tend to stay there, shining steadily. Stars like the sun would brighten somewhat in their first couple of billion years. The brightening of the sun poses a problem called the faint early sun problem: geological evidence indicates that the earth has been warm enough to have liquid water throughout its history, yet the early sun was perhaps 25% less bright than the present sun, and the early earth should have been cold. Perhaps the early earth had a denser atmosphere that trapped heat more than the atmosphere does now. Both geologists and astronomers are actively pursuing research on this question.
Almost every aspect of the life of a main sequence star is determined by one fact: its mass. Very tiny stars emit light feebly and remain cool. Such stars are called red dwarfs. More massive stars are hotter and brighter. Massive stars have more fuel to sustain their output, but their energy output or luminosity is roughly proportional to the fourth power of their mass. A star twice as massive as the sun will emit energy 2 x 2 x 2 x 2 or 16 times the rate of the sun. With only twice as much fuel to sustain it, the star will only shine 2/16 or 1/8 as long as the sun before running out of fuel.
We might compare stars to people. Red dwarf stars spend their energy frugally, like someone with a very limited income, whereas bright, massive blue-white stars run through their fuel quickly, like a lottery winner on a spending spree. Red dwarf stars can last tens or hundreds of billions of years, doling out their fuel at a miserly rate. Our own sun will shine for perhaps 10 billion years, but bright blue-white supergiants like Deneb or Rigel might last only a few million years.
A rough estimate of star lifetimes on the Hertzsprung-Russell diagram is shown below. The hottest stars on the Main sequence don't even last a million years, but red dwarfs last far longer than the age of the Universe. Stars, whether in Hollywood or the sky, follow the principle "Live hard, die young, leave a good-looking corpse."
When main sequence stars run out of available fuel, the nuclear reactions in the center of the star die out. Without the intense radiation pressure produced by these nuclear reactions, the star begins to contract under its own gravity. As matter falls in toward the center of the star, it releases energy that heats the star until finally a new sequence of nuclear reactions begins. The renewed energy output heats the outer part of the star, causing it to expand. At this point the star leaves the main sequence. The outer layers expand and cool, causing the star to redden, but the enormous size of the star gives it a vast surface area through which it emits a tremendous amount of energy. The star becomes a red giant or supergiant.
The most massive stars leave the main sequence soonest. We can see the evolution of stars clearly by plotting H-R diagrams for star clusters, whose members all formed at the same time. Very young clusters may even have remnants of their parent gas clouds still visible and contain nothing but main sequence stars, and even stars that have not yet reached the main sequence. Older clusters have some of their brightest stars leaving the main sequence, and in very old clusters, most of the stars brighter than the sun have left the main sequence. The evolution of the H-R diagram of a star cluster is rather like peeling a strip off a banana; as the star cluster ages, the giant star branch becomes larger and the branching point moves farther down the main sequence.
Since the lifetimes of Main Sequence stars increase rapidly with decreasing brightness, H-R plots of star clusters commonly have a Giant Branch peeling off in classes O, B, and A. Only old clusters will have a Giant Branch peeling off at F, and only the most ancient globular clusters are old enough to have G stars enter the giant phase. The Universe isn't old enough yet for K stars to have gone giant. M and cooler stars are not massive enough to collapse and begin helium fusion, and they will never become giant stars.
All objects in the universe exist because of a balance between gravity and some counteracting force. Four fundamental forces of nature (gravity, electromagnetism, weak nuclear force and strong nuclear force) govern the structure and bonding of atoms. It is fitting that we return to these same basic forces here when we examine matter on the largest scale.
Left to itself, gravity would pull all masses together to a central point. In planets, the counteracting force is the atomic bonding between atoms, and the repulsion between the negatively-charged electrons of neighboring atoms. In normal stars, the counteracting force is the thermal motion of the atoms in the hot gas of the star, and the outward pressure exerted by radiation.
When the radiation pressure in a star falters, gravity begins pulling the gas of the star inward. As the gas falls inward, it gains energy and heats up the interior of the star. Also, the gas becomes more tightly compressed and the pressure increases. Several things can happen, depending on the mass of the star. The temperature and pressure inside the star can rise until a new generation of nuclear reactions begins, the star can continue to collapse until some new counteracting force stops the collapse, or the star can collapse until gravity actually does pull all the mass of the star into a single point: a black hole. The more massive a star is, the more dramatic its end.
Small, faint, red dwarf stars probably never do anything very dramatic. They continue to fuse hydrogen to helium at a miserly rate. Even nearby red dwarf stars are very faint. If there were none within a few dozen light years of earth we would not know they exist at all, but judging from what we see in the space near the sun, red dwarfs are among the most common stars. These stars can continue to shine, if one can use that word for such faint stars, for tens or hundreds of billions of years, gradually cooling as their fuel runs out. Their mass is so small they will never collapse enough to start a new cycle of activity. The atomic repulsion between atoms will counteract gravity. Even then, their surface area is so small they will remain warm for a very long time.
Stars ranging from 10 per cent to several times the mass of the sun go through a different final history. As the star's hydrogen supply begins to run out and its energy output falters, the star begins to contract. As the matter of the star falls inward, it acquires energy and the star heats up. Eventually, the temperature and pressure inside the star get high enough that the helium in the star can begin to fuse to make carbon. The core of the star is very dense, and its energy output heats the outer gases of the star, causing them to expand. The star swells enormously, becoming perhaps as large in diameter as the earth's orbit: 300 million kilometers (186 million miles). A red giant consists of a dense core and a vast but very thin outer atmosphere. It has a huge surface area to radiate energy, so red giants are very luminous, but the energy is spread thinly. The surface temperature of the star is low, which is why red giants are red.
Many red giants are unstable. Instead of swelling to a given size and maintaining it, the stars pulsate and vary in brightness. Some red giants pulsate rhythmically, others flare up in irregular bursts. Red giants include many varieties of variable stars.
Red giant stars also shed matter into space. Many giants shed matter steadily, others violently. In their later life cycles, some pulsating giants eject great shells of matter which form luminous envelopes around the star. Because these gas envelopes look disk-like in a telescope, they are called planetary nebulae.
Some giants are massive enough to start other cycles of nuclear reactions after their helium runs out; they fuse carbon and perhaps even heavier elements, but sooner or later all red giants run out of nuclear fuel. Their thin outer envelopes are ejected into space or gathered up into their core. The core collapses until a new counteracting force comes into play. The electron shells of the atoms are crushed out of existence and the electrons wander between densely-packed atomic nuclei. The forces between these electrons prevent the star from collapsing further. By this time, the star may be only about as large in diameter as the earth, even though it is as massive as the sun. The matter of the star, called degenerate matter, is so dense that a teaspoonful would weigh many tons on earth. This final stage of the star is a white dwarf. White dwarfs are very hot, but their surface area is so small that they are very faint and lose their heat very slowly. The coolest known white dwarf is about 3900 K. The Universe is not old enough for any white dwarf stars to have cooled completely.
In all likelihood, the sun will become a red giant. In about 10 billion years, all the hydrogen in the core of the sun will have been used and the sun will start to contract under its own gravity. It will heat up and brighten as it does, probably making the earth too hot for life. When helium begins to fuse in the sun's core, the outer gases of the sun will expand, probably enveloping the earth. The gas will be hot, but very thin, and for several thousand years the earth may actually orbit within the star, slowly heating up by contact with the thin hot gas, and eventually being destroyed as friction with the gas causes the earth to spiral into the hotter interior of the star. Eventually the sun will eject its outer envelope, or absorb it, leaving only its core as a white dwarf. In the unlikely event the earth survives the red giant stage, the final white dwarf will emit only a fraction of the present energy of the sun and the earth will be frozen solid.
If binary or multiple stars are far apart, they will evolve independently of one another. However, astonishing things happen when multiple stars are close together. When one of the stars reaches the red giant phase, it can swell large enough to exchange gas with its partner star. What happens depends on the ages of the two stars, their masses, and how rapidly they exchange matter.
If the partner star is a white dwarf, some very violent events can happen. White dwarfs have used up all their nuclear fuel. If a large amount of hydrogen falls onto a white dwarf, it can undergo nuclear fusion right on the surface of the star. The resulting outburst will cause the star to brighten briefly by hundreds of times, becoming a nova (Latin, new). Even more dramatic outbursts are possible: the white dwarf can accumulate mass until its internal pressures become great enough for the next generation of nuclear reactions begin. When that happens, the renewed nuclear activity will blast off the outer layers of the white dwarf, creating a type I supernova.
Very massive stars have a more dramatic end yet: they become Type II Supernovae, stars that explode and briefly outshine all the other stars in the galaxy combined. Because old stars lose mass in the red giant stage, it is hard to predict exactly which stars will meet this most violent of fates.
Red giants with 20 or so times the mass of the sun develop cores with a concentric shell structure. Each shell is hotter and denser than the one outside. In the outermost shell, hydrogen fuses to helium as in any normal star. In the next shell in, helium fuses to carbon. Succeeding shells within fuse carbon, oxygen, neon, and silicon until, at the center, there is a core of inert iron. Iron cannot yield energy by fusing to make heavier atoms, so this innermost core is the end product of fusion in the star.
At first glance, it looks as if the star can turn entirely to iron, because each shell uses up the residue from the shell outside it. But there is a catch: eventually the core becomes so massive that it cannot withstand the pressure of gravity. The core collapses until a new counteracting force halts the collapse. The only force capable of halting the collapse is the strong nuclear force. The core of the star collapses until it becomes, in effect, a gigantic atomic nucleus. The core becomes a neutron star. It may have more mass than the sun but be only 15 kilometers (10 miles) across. A teaspoonful of this matter would weigh thousands of tons on earth.
Where there was once the hot, dense core of the star, there is briefly a void. The neutron star core, more massive than the sun but not much bigger than a mountain on earth, has an incredible gravitational pull. At a distance of 10,000 kilometers (6,000 miles) a neutron star as massive as the sun exerts a gravitational pull about 1,300 times stronger than that on the surface of the earth. In this enormous gravitational field, the matter of the star falls inward, reaching perhaps a tenth of the speed of light.
The results are, to say the least, impressive. Several times the mass of our sun crashes into the surface of a neutron star at up to a tenth of the speed of light. This gas is heated and compressed beyond anything we can imagine. We can do the calculations and write the numbers, but nobody can really comprehend the titanic amount of energy involved. Nuclear reactions run rampant; there is enough energy available and a high enough density of fast-moving particles to build nuclei as heavy as plutonium and probably far heavier. The neutron star core itself is almost incompressible even under these conditions, and the impacting matter rebounds as a shock wave. This event, called the core bounce, tears the star apart. In a few hours, the shock wave reaches the surface, tearing off the outer layers and exposing the hot interior of the star. The star brightens to billions of times its normal brightness.
Supernovae are commonly seen in distant galaxies. If our own galaxy has supernovae as often as other galaxies, there is probably one every few years. Yet only a dozen or so supernovae in our own galaxy have been witnessed from earth. The rest have been obscured by gas and dust clouds in our galaxy.
The supernova of 1572 was brighter than Venus and could be seen in broad daylight, even though it was 5000 light years away. Another supernova occurred in 1604. These events came at a critical time, just as astronomers were beginning to question the ancient notion that the heavens were perfect and unchanging. A supernova in 1006 rivaled the moon in brightness. Most supernovae have absolute magnitudes of about -21: at a distance of 32.6 light years it would shine at magnitude -21 and far outshine the moon. In 1885, a supernova in the Andromeda galaxy reached magnitude 7. Across 2.2 million light years, that single star was almost bright enough to see with the unaided eye. All these events, however, happened before astronomers had the observing techniques to study supernovae in detail.
Almost four centuries of not-very-patient waiting ended in 1987 when the first supernova since 1604 visible to the unaided eye appeared. The supernova occurred not in our galaxy, but in the Small Magellanic Cloud, a small satellite galaxy of our own about 180,000 light years away.
Almost from the beginning, supernova 1987 followed a script all its own. It only reached magnitude 3, not the magnitude 1 that astronomers expected, but it stayed at peak brightness much longer than most supernovae. Most paradoxical of all, the star that produced the supernova was not an aging red supergiant, but a seemingly stable blue-white supergiant. It now appears that some red supergiants can lose their outer envelopes quietly, revealing their hot interiors more clearly.
What might it be like on a planet orbiting a star that went supernova? Imagine the planet receives as much radiation as we get from the sun. It is hard enough to imagine the energy output of the sun, let alone a supernova, so let us scale things down a bit first by asking what it would take to match the sun's output at a distance of only one kilometer. The sun emits 77 megatons of energy every second, but it is 150 million kilometers away, and the intensity of radiation drops off as the square of the distance. To find out what energy output we need for a distance of one kilometer, we must divide the sun's output by the square of its distance. The answer comes out to about .0000035 megatons, or the equivalent of 3.5 tons of high explosive. It is not hard to picture an explosion of 3.5 tons of high explosive (a truckload) a kilometer away giving off a one-second burst of heat and light that rivals the sun.
Now, if our hypothetical star were to go supernova, its brightness would increase 100 billion times, or be equivalent to 350 billion tons of explosive a kilometer away. 350 billion tons translates to 350,000 megatons, or much more than all the energy in all the earth's nuclear weapons. In other words, the planet would receive a blast of heat and light equivalent to having every nuclear weapon on earth detonated at the same time a kilometer away -- and this intensity would last for days. It is no exaggeration to say that the planet would be vaporized.
Fortunately, the melancholy idea of a star going supernova and frying the life on its planets is unlikely. The very massive stars that produce supernovae do not shine long enough for life to evolve beyond the simplest forms, and many may not even last long enough for planets to finish the accretion process.
For many thousands of years, a supernova can be recognized by the expanding shell of gas blasted off the star. Such a shell is called a supernova remnant. The former star itself is often detectable as a pulsating radio source, or pulsar. Pulsars emit tremendous bursts of energy in radio, visible light, and x-ray wavelengths. These bursts appear to originate in small regions of the neutron star, perhaps due to matter falling onto the neutron star. These pulses are extremely regular, ranging in period from .001 second to a few seconds. The period of the pulses is the rotation period of the neutron star. As the parent star collapses, its rotation speeds up, just as a pirouetting skater speeds up by drawing in her arms. If the sun, with a diameter of 864,000 miles (1.3 million km) were to shrink to a netron star 10 miles (16 km) in diameter, it would shrink to 1/86400 of its present size, and its rotation would speed up 86400 times. Instead of rotating every 28 days, the neutron sun would rotate in 28 seconds.
Supernova remnants often appear to be associated with areas of new star formation, and it is very likely that a supernova triggers the formation of new stars. If the expanding blast wave from a supernova strikes a cloud of interstellar gas and dust, it can compress the cloud enough that parts of the cloud start to contract gravitationally. The supernova debris will also mingle with the cloud, enriching it in heavy elements.
In normal stars, it is not possible to form large amounts of elements much heavier than iron. Elements like lead, gold, and uranium can form only in supernovae, as far as we can tell. The fact that we find these elements in the solar system is a sign that our sun is perhaps a third-generation star. The sun is only a third as old as the milky way galaxy, and there was time for many cycles of stellar birth and death before the sun formed. The implications of this idea are profound. Every atom in us formed in a star billions of years ago and many light-years away.
If the collapsed core of a supernova is more than 1.4 times as massive as the sun, it will not form a neutron star. Instead, there is no known force that can halt the gravitational collapse. The star will contract until its gravity is so immense that not even light can escape. As far as we know, the star will contract until it becomes a point, detectable only by its gravity. What happens to the matter in the star? We can only speculate, but it is possible, according to some theories in physics, that the matter might re-emerge somewhere else in space and time.
These bizarre objects, called black holes, are favorite topics of speculation among science-fiction and popular magazine articles, but have any actually been detected? Possibly some have. A black hole orbiting another star might draw matter from the companion star. As the gas fell into the black hole, it would accelerate to enormous velocities and emit intense x-rays. There are a few x-ray sources that are so massive that they appear likely to be black holes.
Where did atoms come from, and how did there come to be so many different kinds? We can answer this question in quite a bit of detail, thanks to what we know about nuclear reactions in particle accelerators, reactors, and nuclear bombs.
The sun gets its energy by fusing four hydrogen nuclei (protons) to make a helium nucleus. This process is actually fairly complicated and proceeds in several steps. The final nucleus contains four particles, or nucleons: two protons and two neutrons.
Beyond helium we encounter a bottleneck. If we try to add a neutron or proton to a helium nucleus, the new nucleus disintegrates almost instantly. There is no nucleus with five nucleons that lasts long enough to be a basis for heavier elements. Perhaps we can fuse two helium nuclei together? But it turns out that there are no long-lasting nuclei with eight nucleons either.
However, there is a way around this bottleneck. When stars collapse to form red giants, the temperatures and pressures inside the star become high enough for three-way collisions of helium nuclei to occur. These collisions are rare, to be sure, but they occur often enough to form heavier elements in the quantities we observe. The product of these collisions has 12 nucleons: 6 protons and 6 neutrons, and is a carbon nucleus.
What about lithium, beryllium, and boron, the elements between helium and carbon? These atoms form only when collisions knock nucleons off of heavier nuclei, a process called spallation, and they tend to be destroyed in the interiors of stars. Compared to carbon, they are rare in the universe. Actually, it is probably a good thing that it is so hard to form heavy atoms. If it were easier, stars might long ago have fused all their hydrogen to heavy atoms and there would be no available energy left in the universe.
Once carbon forms, it serves as a base for building heavier atoms. Nucleons can be added one-by-one, or by fusing more helium nuclei to an existing nucleus. Building elements by fusing helium nuclei is a common process, and elements with even numbers of protons in their nuclei are more abundant than atoms with odd numbers. But each heavier nucleus takes more energy to make and gives less energy back, until at iron, with 26 protons, the process ends. Beyond iron, it takes more energy to make a nucleus than the nuclear reaction gives back. Random collisions build some heavier nuclei, but the abundance of elements drops off sharply after iron.
Very heavy atoms like gold are extremely rare in the universe, and seem to form only in two ways. One way, the s-process (for slow), is the addition of neutrons one by one to existing nuclei. Eventually, once enough neutrons have been added, one of them breaks down to a proton plus an electron. The electron is expelled, a process called beta-decay, and the atomic number of the element increases by one. Since the process is slow, many short-lived radioactive nuclei decay away before another neutron can be added to the nucleus. The atoms that can form by the s-process are limited to a narrow band of stable and long-lived radioactive nuclei, and the r-process ends with the formation of lead. The r-process happens in certain red giant stars, and the heavy elements are blown into space as the star loses its outer shells.
the fierce environment in the core of a supernova. Much of the light given off by a supernova turns out to be due to the decay of radiocative nickel. In a supernova, there is so much energy available that particles can pile onto nuclei at a tremendous rate. Atoms at least as heavy as plutonium form this way, and probably atoms far heavier than that. These atoms require more energy to form than our most powerful particle accelerators can produce, but nuclear physics predicts that nuclei with about 110 protons might have very long lifetimes, possibly long enough to be left over from the formation of the earth. Some physicists have attempted to find such super-heavy elements in rocks, so far without success.
Created 26 March 1998, Last Update 30 March 1998
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