The life cycles of stars follow three general patterns, each associated with a range of initial mass.
We will talk about intermediate mass stars.
A star develops from giant molecular cloud (also called dark nebula).These molecular clouds are clumpy with varying density. Their density varies from a few molecules to millions of molecules per cubic centimeter. These molecules are mainly of hydrogen. Stars form from the densest regions of molecular cloud, which are called cloud cores.
Due to the force of its own gravity, the cloud begins to collapse inward, thereby becoming smaller. As the cloud shrinks, it rotates more and more rapidly. The outermost parts of the cloud form a spinning disk. The inner parts become a roughly spherical clump, which continues to collapse.
The collapsing material becomes warmer, and its pressure increases. But the pressure tends to counteract the gravitational force that is responsible for the collapse. Eventually, therefore, the collapse slows to a gradual contraction. The inner parts of the clump form a protostar. Surrounding the protostar is an irregular sphere of gas and dust that had been the outer parts of the clump.
A cloud that eventually develops into an intermediate-mass star takes about 100,000 years to collapse into a protostar. As a protostar, it has a surface temperature of about 4000 K. It may be anywhere from a few times to a few thousand times as luminous as the sun, depending on its mass.
When the temperature and pressure in the protostar's core become high enough, nuclear fusion begins. Nuclear fusion is a joining of two atomic nuclei to produce a larger nucleus.
Nuclei that fuse are actually the cores of atoms. A complete atom has an outer shell of one or more particles called electrons, which carry a negative electric charge. Deep inside the atom is the nucleus, which contains almost all the atom's mass. The simplest nucleus, that of the most common form of hydrogen, consists of a single particle known as a proton. A proton carries a positive electric charge. All other nuclei have one or more protons and one or more neutrons. A neutron carries no net charge, and so a nucleus is electrically positive. But a complete atom has as many electrons as protons. The net electric charge of a complete atom is therefore zero.
However, under the enormous temperatures and pressures near the core of a protostar, atoms lose electrons. The resulting atoms are known as ions, and the mixture of the free electrons and ions is called a plasma. Atoms in the core of the protostar lose all their electrons, and the resulting bare nuclei approach one another at tremendous speeds. Under ordinary circumstances, objects that carry like charges repel each other. However, if the core temperature and pressure become high enough, the repulsion between nuclei can be overcome and the nuclei can fuse.
When two relatively light nuclei fuse, a small amount of their mass turns into energy. Thus, the new nucleus has slightly less mass than the sum of the masses of the original nuclei. Albert Einstein discovered the relationship E = mc2 that indicates how much energy is released when fusion occurs.
The speed of light is 186,282 miles (299,792 kilometers) per second. This is such a large number that the conversion of a tiny quantity of mass produces a tremendous amount of energy. For example, complete conversion of 1 gram of mass releases 90 trillion joules of energy. This amount of energy is roughly equal to the quantity released in the explosion of 22,000 tons (20,000 metric tons) of TNT. This is much more energy than was released by the atomic bomb that the United States dropped on Hiroshima, Japan, in 1945 during World War II. The energy of the bomb was equivalent to the explosion of 13,000 tons (12,000 metric tons) of TNT.
In the core of a protostar, fusion begins when the temperature reaches about 1 million K. This initial fusion destroys nuclei of certain light elements. These include lithium 7 nuclei, which consist of three protons and four neutrons. In the process involving lithium 7, a hydrogen nucleus combines with a lithium 7 nucleus, which then splits into two parts. Each part consists of a nucleus of helium 4, two protons and two neutrons. A helium 4 nucleus is also known as an alpha particle.
After the light nuclei are destroyed, the protostar continues to contract. Eventually, the core temperature reaches about 10 million K, and hydrogen fusion begins. The protostar is now a star. In hydrogen fusion, four hydrogen nuclei fuse to form a helium 4 nucleus. There are two general forms of this reaction:
The p-p reaction can occur in several ways, including the following four-step process:
The CNO cycle differs from the p-p reaction mainly in that it involves carbon 12 nuclei. These nuclei consist of six protons and six neutrons. During the cycle, they change into nuclei of nitrogen 15 (7 protons and 8 neutrons) and oxygen 15 (8 protons and 7 neutrons). But they change back to carbon 12 nuclei by the end of the cycle.
Helium nuclei can fuse to form carbon 12 nuclei. However, the core temperature must rise to about 100 million K for this process to occur. This high temperature is necessary because the helium nuclei must overcome a much higher repulsive force than the force between two protons. Each helium nucleus has two protons, so the repulsive force is four times as high as the force between two protons.
The fusion of helium is called the triple-alpha process because it combines three alpha particles to create a carbon 12 nucleus. Helium fusion also produces nuclei of oxygen 16 (8 protons and 8 neutrons) and neon 20 (10 protons and 10 neutrons).
At core temperatures of about 600 million K, carbon 12 can fuse to form sodium 23 (11 protons, 12 neutrons), magnesium 24 (12 protons, 12 neutrons), and more neon 20. However, not all stars can reach these temperatures.
As fusion processes produce heavier and heavier elements, the temperature necessary for further processes increases. At about 1 billion K, oxygen 16 nuclei can fuse, producing silicon 28 (14 protons, 14 neutrons), phosphorus 31 (15 protons, 16 neutrons), and sulfur 32 (16 protons, 16 neutrons).
Fusion can produce energy only as long as the new nuclei have less mass than the sum of the masses of the original nuclei. Energy production continues until nuclei of iron 56 (26 protons, 30 neutrons) begin to combine with other nuclei. When this happens, the new nuclei have slightly more mass than the original nuclei. This process therefore uses energy, rather than producing it.
When hydrogen fusion begins, the protostar is still surrounded by an irregular mass of gas and dust. But the energy produced by hydrogen fusion pushes away this material as a protostellar wind. In many cases, the disk that is left over from the collapse channels the wind into two narrow cones or jets. One jet emerges from each side of the disk at a right angle to the plane of the disk. The protostar has become a T-Tauri star, a type of object named after the star T in the constellation Taurus (the Bull). A T-Tauri star is a variable star, one that varies in brightness.
The T-Tauri star contracts for about 10 million years. It stops contracting when its tendency to expand due to the energy produced by fusion in its core balances its tendency to contract due to gravity. By this time, hydrogen fusion in the core is supplying all the star's energy. The star has begun the longest part of its life as a producer of energy from hydrogen fusion, the main-sequence phase. The name of this phase comes from a part of the H-R diagram.
Any star -- whatever its mass -- that gets all its energy from hydrogen fusion in its core is said to be "on the main sequence" or "a main-sequence star." The amount of time a star spends there depends on its mass. The greater a star's mass, the more rapidly the hydrogen in its core is used up, and therefore the shorter is its stay on the main sequence. An intermediate-mass star remains on the main sequence for billions of years.
When all the hydrogen in the core of an intermediate-mass star has fused into helium, the star changes rapidly. Because the core no longer produces fusion energy, gravity immediately crushes matter down upon it. The resulting compression quickly heats the core and the region around it. The temperature becomes so high that hydrogen fusion begins in a thin shell surrounding the core. This fusion produces even more energy than had been produced by hydrogen fusion in the core. The extra energy pushes against the star's outer layers, and so the star expands enormously.
As the star expands, its outer layers become cooler, so the star becomes redder. And because the star's surface area expands greatly, the star also becomes brighter. The star is now a red giant.
Eventually, the core temperature reaches 100 million K, high enough to support the triple-alpha process. This process begins so rapidly that its onset is known as helium flash.
As the triple-alpha process continues, the core expands, but its temperature drops. This decrease in temperature causes the temperature of the hydrogen-burning shell to drop. Consequently, the energy output of the shell decreases, and the outer layers of the star contract. The star becomes hotter but smaller and fainter than it had been as a red giant. This change occurs over a period of about 100 million years.
At the end of this period, the star is in its horizontal branch phase, named for the position of the point representing the star on the H-R diagram. The star steadily burns helium and hydrogen, and so its temperature, size, and luminosity do not change significantly. This phase lasts for about 10 million years.
When all the helium in the core has fused, the core contracts and therefore becomes hotter. The triple-alpha process begins in a shell surrounding the core, and hydrogen fusion continues in a shell surrounding that. Due to the increased energy produced by the burning in the shells, the star's outer layers expand. The star becomes a giant again, but it is bluer and brighter than it was the first time.
On the H-R diagram, the point representing the star has moved upward and to the right along a line known as the asymptotic (as ihm TOT ihk) giant branch (AGB). The star is therefore called an AGB star.
An AGB star's core is so hot and its gravitational grip on its outermost layers is so weak that those layers blow away in a stellar wind. As each layer blows away, a hotter layer is exposed. Thus, the stellar wind becomes even stronger. Out in space, a succession of new, fast winds slam into old, slow winds that are still moving away from the star. The collisions produce dense shells of gas, some of which cool to form dust.
In just a few thousand years, all but the hot core of an AGB star blows away, and fusion ceases in the core. The core illuminates the surrounding shells. Such shells looked like planets through the crude telescopes of astronomers who studied them in the 1800's. As a result, the astronomers called the shells planetary nebulae -- and today's astronomers still do. The word nebulae is Latin for clouds.
After a planetary nebula fades from view, the remaining core is known as a white dwarf star. This kind of star consists mostly of carbon and oxygen. Its initial temperature is about 100,000 K.
Because a white dwarf star has no fuel remaining for fusion, it becomes cooler and cooler. Over billions of years, it cools more and more slowly. Eventually, it becomes a black dwarf -- an object too faint to detect. A black dwarf represents the end of the life cycle of an intermediate-mass star.
High-mass stars, those with more than 8 solar masses, form quickly and have short lives. A high-mass star forms from a protostar in about 10,000 to 100,000 years.
High-mass stars on the main sequence are hot and blue. They are 1,000 to 1 million times as luminous as the sun, and their radii are about 10 times the solar radius. High-mass stars are much less common than intermediate- and low-mass stars. Because they are so bright, however, high-mass stars are visible from great distances, and so many are known.
A high-mass star has a strong stellar wind. A star of 30 solar masses can lose 24 solar masses by stellar wind before its core runs out of hydrogen and it leaves the main sequence.
As a high-mass star leaves the main sequence, hydrogen begins to fuse in a shell outside its core. As a result, its radius increases to about 100 times that of the sun. However, its luminosity decreases slightly. Because the star is now emitting almost the same amount of energy from a much larger surface, the temperature of the surface decreases. The star therefore becomes redder.
As the star evolves, its core heats up to 100 million K, enough to start the triple-alpha process. After about 1 million years, helium fusion ends in the core but begins in a shell outside the core. And, as in an intermediate-mass star, hydrogen fuses in a shell outside that. The high-mass star becomes a bright red supergiant.
When the contracting core becomes sufficiently hot, carbon fuses, producing neon, sodium, and magnesium. This phase lasts only about 10,000 years. A succession of fusion processes then occur in the core. Each successive process involves a different element and takes less time. Whenever a different element begins to fuse in the core, the element that had been fusing there continues to fuse in a shell outside the core. In addition, all the elements that had been fusing in shells continue to do so. Neon fuses to produce oxygen and magnesium, a process that lasts about 12 years. Oxygen then fuses, producing silicon and sulfur for about 4 years. Finally, silicon fuses to make iron, taking about a week.
At this time, the radius of the iron core is about 1,900 miles (3,000 kilometers). Because further fusion would consume energy, the star is now doomed. It cannot produce any more fusion energy to balance the force of gravity.
When the mass of the iron core reaches 1.4 solar masses, violent events occur. The force of gravity within the core causes the core to collapse. As a result, the core temperature rises to nearly 10 billion K. At this temperature, the iron nuclei break down into lighter nuclei and eventually into individual protons and neutrons. As the collapse continues, protons combine with electrons, producing neutrons and neutrinos. The neutrinos carry away about 99 percent of the energy produced by the crushing of the core.
Now, the core consists of a collapsing ball of neutrons. When the radius of the ball shrinks to about 6 miles (10 kilometers), the ball rebounds like a solid rubber ball that has been squeezed.
All the events from the beginning of the collapse of the core to the rebounding of the neutrons occur in about one second. But more violence is in store. The rebounding of the ball of neutrons sends a spherical shock wave outward through the star. Much of the energy of the wave causes fusion to occur in overlying layers, creating new elements. As the wave reaches the star's surface, it boosts temperatures to 200,000 K. As a result, the star explodes, hurling matter into space at speeds of about 9,000 to 25,000 miles (15,000 to 40,000 kilometers) per second. The brilliant explosion is known as a Type II supernova.
Supernovae enrich the clouds of gas and dust from which new stars eventually form. This enrichment process has been going on since the first supernovae billions of years ago. Supernovae in the first generation of stars enriched the clouds with materials that later went into making newer stars.
Three generations of stars may exist. Astronomers have not found any of what would be the oldest generation, Population III, stars. But they have found members of the other two generations. Population II stars, which would be the second generation, contain relatively small amounts of heavy elements. The more massive ones aged and died quickly, thereby contributing more nuclei of heavy elements to the clouds. For this reason, Population I stars, the third generation, contain the largest amounts of heavy elements. Yet these quantities are tiny compared with the amount of hydrogen and helium in Population I stars. For example, elements other than hydrogen and helium make up from 1 to 2 percent of the mass of the sun, a Population I star.
After a Type II supernova blast occurs, the stellar core remains behind. If the core has less than about 3 solar masses, it becomes a neutron star. This object consists almost entirely of neutrons. It packs at least 1.4 solar masses into a sphere with a radius of about 6 to 10 miles (10 to 15 kilometers).
Neutron stars have initial temperatures of 10 million K, but they are so small that their visible light is difficult to detect. However, astronomers have detected pulses of radio energy from neutron stars, sometimes at a rate of almost 1,000 pulses per second.
A neutron star actually emits two continuous beams of radio energy. The beams flow away from the star in opposite directions. As the star rotates, the beams sweep around in space like searchlight beams. If one of the beams periodically sweeps over Earth, a radio telescope can detect it as a series of pulses. The telescope detects one pulse for each revolution of the star. A star that is detected in this way is known as a pulsar.
If the stellar core remaining after the supernova explosion has about 3 or more solar masses, no known force can support it against its own gravitation. The core collapses to form a black hole, a region of space whose gravitational force is so strong that nothing can escape from it. A black hole is invisible because it traps even light. All its matter is located at a single point in its center. This point, known as a singularity, is much smaller than an atomic nucleus.
Low-mass stars, ranging from 0.1 to 0.5 solar mass, have surface temperatures less than about 4,000 K. Their luminosities are less than 2 percent of the solar luminosity. Low-mass stars use hydrogen fuel so slowly that they may shine as main-sequence stars for 100 billion to 1 trillion years. This life span is longer than the present age of the universe, believed to be 10 billion to 20 billion years. Therefore, no low-mass star has ever died. Nevertheless, astronomers have determined that low-mass stars will never fuse anything but hydrogen. Thus, as these stars die, they will not pass through a red-giant phase. Instead, they will merely cool to become white dwarfs, then black dwarfs.
Binary stars develop from two protostars that form near each other. More than 50 percent of what seem to the unaided eye to be single stars are actually binaries.
One star in a binary system can affect the life cycle of the other if the two stars are sufficiently close together. Between the stars is a location called the Lagrange point, named for the French mathematician Joseph Louis Lagrange, where the star's gravitational forces are exactly equal. If one of the stars expands so much that its outer layers pass the Lagrange point, the other star will begin to strip away those layers and accumulate them on its surface.
This process, called mass transfer, can take many forms. Mass transfer from a red giant onto a main-sequence companion can add absorption lines of carbon or other elements to the spectrum of the main- sequence star. But if the stars are close together, the material will flow in the opposite direction when the giant star becomes a white dwarf. The matter will spiral in toward the dwarf, forming a hot disk around it. The disk will flare brilliantly in visible and ultraviolet radiation.
If the giant star leaves behind a neutron star or a black hole instead of a white dwarf, an X-ray binary may form. In this case, the matter transferred from the main-sequence star will become extremely hot. When this matter strikes the surface of the neutron star or is pulled into the black hole, it will emit X rays.
In a third case, the red giant becomes a white dwarf, and the main-sequence star becomes a red giant. When enough gas from the giant accumulates on the dwarf's surface, gas nuclei will fuse violently in a flash called a nova. In some cases, so much gas will accumulate that its weight will cause the dwarf to collapse. Almost instantly, the dwarf's carbon will fuse, and the entire dwarf will explode in a Type I supernova. This kind of explosion is so bright that it can outshine an entire galaxy for a few months.
