Stellar evolution

Stars are “born” in interstellar clouds of gas and dust called nebulae. They generate energy by various mechanisms for most of their “lives.” They “die” when they finally run out of ways to produce any more energy.

Overview

One of the historical goals of astronomers has been to determine how stars are born, how they live, and how they die. This community now believes a comprehensive understanding now exists. Stars go through a series of changes that are referred to as stellar evolution or stellar life cycles, in analogy to the life stages of living organisms.

Stars are born in nebulae, large clouds of gas and dust in interstellar space. The clouds typically are several tens of light-years in diameter and contain enough matter to form hundreds to thousands of stars. The gas is mostly hydrogen (about 71 percent by mass), with some helium (about 27 percent by mass) and small amounts of heavier elements. The dust particles are small, solid grains of carbon, Silicateminerals, iron compounds, and ice, probably about 0.001 millimeters in size, on average. The nebulae in which stars form are cold, with temperatures on the order of 10 kelvins. At such low temperatures, atoms can bond together to form molecules, so these nebulae are referred to as molecular clouds. Also as a result of the low temperatures, the gas pressure is very weak and can barely keep the cloud from contracting from its self-gravity. In denser parts of the cloud, self-gravity may overcome the weak pressure and start that part of the cloud collapses. The trigger for this increased density may be an encounter with a Spiral arm density wave as the Nebula orbits the center of its galaxy, strong stellar winds from a nearby star that already formed, or shock waves from a nearby Supernova explosion.

A molecular cloud usually breaks up into separate clumps that collapse on their own. As each clump collapses, gravitational energy is rapidly converted into thermal energy, heating the gas and causing it to begin feebly emitting radio waves. When a sufficiently dense core has formed at the center of the clump, the contraction slows and the object is called a protostar. Typically, many neighboring protostars form at about the same time. A protostar continues to contract because of gravity. It grows hot enough to shine in the infrared part of the electromagnetic spectrum. What little visible light the protostar emits is blocked by the shroud of dust that surrounds it. Eventually the dust shroud dissipates, some of it joining the growing equatorial disk around the protostar (in which planets may ultimately form) and some of it blown away as the protostar becomes hotter and brighter. As the protostar continues to contract and get hotter, collisions break molecules apart into individual atoms and electrons are stripped off the atoms, ionizing the gas.

When the temperature at the center of the contracting protostar reaches several millions of degrees (in kelvins), hydrogen nuclei start to fuse together to form helium nuclei. The energy released by this nuclear fusion reaction stops the contraction, and the protostar becomes a main sequence star. It takes higher-mass protostars less time to reach the main sequence stage than lower-mass protostars, because their greater mass means stronger gravity and more rapid contraction. The protostar stage lasts on the order of a hundred thousand years for a star with 10 solar masses, a few tens of millions of years for a star like the Sun, with 1 solar mass, and from several hundred million to a billion years for a star with 0.1 solar mass.

In hydrogen fusion, four hydrogen nuclei fuse together to form one helium nucleus. The mass of the four hydrogen nuclei is slightly greater than the mass of the one helium Nucleus produced, and the excess mass is converted into energy according to Albert Einstein’s formula E = mc², which states energy E and mass m are equivalent and are related by a physical constantc (the speed of light squared). To overcome the electrical repulsion that the positively charged hydrogen nuclei (bare protons) have for each other requires high temperatures (so the nuclei are moving quickly) and high densities (so the nuclei are close together). These are the conditions in the cores of main sequence stars.

A main sequence star is in a state of hydrostatic equilibrium, which means the self-gravity trying to make the star contract is balanced by the pressure trying to make the star expand. The energy released by hydrogen fusion in the core provides the energy that the star radiates into space. The main sequence stage is the longest and most stable part of a star’s energy-producing life. A star remains a main sequence star as long as it has hydrogen in its core to fuse to helium. Massive main sequence stars have more fuel, but they consume that fuel much more rapidly, so their main sequence lives are relatively short. For example, a 30-solar-mass star has a main sequence lifetime of about 5 million years. The Sun, with 1 solar mass, has a main sequence lifetime of about 10 billion years. (Since the Sun is currently about 4.5 billion years old, it has progressed about halfway through its main sequence life.) Low-mass main sequence stars have less fuel, but they consume it much more slowly, so their main sequence lives are much longer, from about 30 billion years for a star with 0.5 Solar mass to as much as a trillion years for a star with 0.1 solar mass. These main sequence lifetimes are longer than the age of the universe, so every low-mass main sequence star that ever formed is still a low-mass main sequence star.

When all the hydrogen in the core has been fused into helium, the core contracts and heats up. Hydrogen fusion is transferred to a still hydrogen-rich shell surrounding the contracting helium-rich core. This causes the outer layers of the star to expand and cool, and the expanding surface of the star becomes red in color. The star ceases to be a main sequence star and becomes a red giant or red supergiant. Stars similar to the Sun, with about 1 or 2 solar masses, become red giants, ten to one hundred times bigger than the present Sun. Massive stars become red supergiants, one hundred to one thousand times bigger than the present Sun.

When the temperature of the contracting helium-rich core reaches about 100 million kelvins, helium fusion is ignited. Three helium nuclei fuse to form one carbon nucleus. Add another helium nucleus, and an oxygen nucleus forms. The total mass of the three or four helium nuclei is a bit greater than the mass of the carbon or oxygen nucleus, and again this mass excess gets converted into energy. This fusion reaction supplies the star with energy for much less time than did hydrogen fusion when it was a main sequence star.

When the helium in the core is exhausted, the core once again contracts and heats up. A star like the Sun is not massive enough for its core to shrink enough to get hot enough to start further nuclear fusion reactions. Thermal pulsations and a strong Stellar wind blow off its outer layers in one or more shells of gas. The expanding shell of gas surrounding what is left of the star is called a planetary nebula. Planetary nebulae have no association with planets. The name originated in the 1800s because in the telescopes then in use, the objects looked round, like planets, and fuzzy, like nebulae. The central star of a planetary nebula is the exposed former core of the star. It contracts as much as it can and becomes a white dwarf star composed of carbon and maybe oxygen. White dwarf stars have about the mass of the Sun packed into a sphere about the size of the Earth. This makes them very dense, averaging about 1 metric ton per cubic centimeter. White dwarf stars can no longer generate energy; they cannot contract to release gravitational energy, and thus they cannot get any hotter to be able to tap other nuclear fusion reactions. They shine only because they are very hot. As they shine, they radiate their energy away and cool, gradually becoming black dwarfs. These can be described as cold, dark, dead stellar “embers."

The core of a massive star (more than about eight solar masses) can shrink and get hot enough to go through a series of nuclear fusion reactions, one after another, to produce heavier atomic nuclei. This stops with the production of iron nuclei, since iron is the heaviest nucleus that can form through fusion reactions that release energy. To form heavier nuclei through fusion requires the input of energy. The iron core collapses, and the outer layers collapse on top of it and rebound, sending shock waves through the star. This tears the star apart in a Type II supernova explosion. In a few minutes, it releases more energy than it produced by nuclear fusion reactions during its entire preceding life. A Type II supernova becomes about a billion times more luminous than the Sun before it gradually fades away. Some of the prodigious energy released in the explosion goes into fusion reactions forming elements heavier than iron. Much of the star’s matter is violently ejected into interstellar space at speeds ranging from thousands to a few tens of thousands of kilometers per second, thereby enriching the interstellar material in elements heavier than helium.

If the mass of the stellar remnant that remains after the Type II supernova explosion is less than approximately 2 to 3 solar masses, it becomes a neutron star. A Neutron star has a radius of only 15 kilometers and a density of a billion metric tons per cubic centimeter. If the remaining mass of the Supernova remnant is greater than about 3 solar masses, however, no known force can stop its collapse to a black hole. The Gravitational field of a Black hole is so great that nothing, not even light, can escape from it. Consequently, a black hole cannot be directly “seen” in any part of the electromagnetic spectrum. A black hole can be detected only by its effects, primarily through its gravity, on nearby objects.

Knowledge Gained

The Earth and all life on it exist only because of the life and death of massive stars. Our current understanding of the Big Bang (the primordial “explosion” that created the universe about 13 to 14 billion years ago) is that it produced only hydrogen and helium with very small traces of Lithium and beryllium. The first stars that formed in the earliest days of our galaxy consisted only of hydrogen and helium. Gas planets like Jupiter might have formed around those first stars, but not the rocky, metallic planets like Earth. Carbon-based life could never have developed. The atoms of all the elements heavier than helium that make up our bodies and the Earth we live on were produced by nuclear fusion reactions in massive stars before and during their deaths as Type II supernovae. These heavy elements were spewed out into interstellar space by the supernova explosions, where they enriched the interstellar clouds of gas from which new stars formed. By the time the Sun and the solar system began to form in one of these nebulae about 4.5 billion years ago, about 2 percent of its mass consisted of elements heavier than helium. This provided the material for rocky/metallic planets to form in the inner solar system and for carbon-based life-forms to develop on one of them, our Earth. The production and dispersal of heavy elements by massive stars, with the resulting enrichment of nebulae, continues to the present time, so that today about 5 or 6 percent of interstellar matter consists of elements heavier than helium.

When the Earth’s solar system was formed, again 4.5 billion years in the past, it originated from a dense cloud of gas and dust. This material later evolved into a spinning solar nebula. At its center, gravity accumulated an increasing amount of matter, most importantly hydrogen and helium. The combination of these two elements began to transmit enormous quantities of energy which resulted in the formation of the sun.

Our understanding of the evolution of stars like the Sun reveals the future of our solar system, including Earth. About 5 billion years from now, the Sun will run out of hydrogen in its core. It will expand until its radius is about one hundred times larger than it is now, swallowing the planet Mercury and possibly Venus. At that point it will be a red giant, about a thousand times more luminous than it is now. Temperatures on Earth will reach between 1,000 and 2,000 kelvins; our atmosphere will escape into space, our oceans will boil away, surface rocks will at least partly melt, and life will not survive. Eventually, strong Solar wind and thermal pulses will blow the Sun’s outer envelope back into space, producing a planetary nebula. What is left of the Sun will become a white dwarf, with perhaps one-half to two-thirds of its present mass. Initially, this white dwarf Sun will have a surface temperature of about 100,000 kelvins, but since it will be unable to generate more energy, it will cool, rapidly at first, then ever more slowly, until it becomes a cold, dark, dead black dwarf.

Context

New and improved instrumentation, including detectors in space above Earth’s atmosphere, made it possible to observe the solar system and deep space with sharper Resolution and in all parts of the electromagnetic spectrum. Advances in theoretical physics and the widespread use of computers led to the construction of more and better models of the stellar interiors and study of how they change with time.

In the late 1890s and early 1900s, George Ellery Hale oversaw the design and construction of several large telescopes—the 40-inch refractor at Yerkes Observatory in Wisconsin, and the 60-inch and 100-inch reflectors on Mount Wilson in California. For the next half-century, they were the largest telescopes in the world, and their use led to many observational discoveries in all branches of astronomy.

The early 1900s also saw the development of the Hertzsprung-Russell (or H-R) diagram, a graph used for plotting Stellar luminosity versus surface temperature. It turns out that the location of a star in this diagram indicates many things about the star; besides its luminosity and surface temperature, these include its radius, in some cases its approximate mass, and its evolutionary stage. Consequently, the H-R diagram has proved to be a powerful tool for tracing the life cycles of stars of various masses.

In 1926, Sir Arthur Stanley Eddington published The Internal Constitution of the Stars, a book that established much of the formalism still used today to construct models of stellar interiors. Its most serious deficiencies, of which Eddington was well aware, involved the sources of stellar opacities and energy generation. In the 1930s, Subrahmanyan Chandrasekhar used relativity theory to work out the structure of white dwarf stars, which, since their discovery in the middle to late nineteenth century, had at best been a stellar enigma and which many astronomers considered to be a physical impossibility.

Continuing through the 1930s and into the 1940s, the problem of stellar energy generation began to be addressed by advances in nuclear physics, which led to the idea that stars generate energy through nuclear fusion reactions; a few specific fusion reactions were suggested for main sequence stars. Work in nuclear physics also suggested the existence of neutron stars, though many thought that even if they did exist, they could never be detected. (It was not until 1967 that the discovery of pulsars, identified as rapidly rotating, highly magnetic neutron stars, confirmed the existence of neutron stars.)

The late 1950s saw the publication of a number of seminal articles and books. A long article by Margaret Burbidge, Geoffrey Burbidge, William A. Fowler, and Fred Hoyle titled “Synthesis of the Elements in Stars” (1957) traced the various processes by which heavy elements could be produced by Nucleosynthesis in stars. Martin Schwarzschild’s book Structure and Evolution of the Stars (1958) described in an accessible way the methods for computing models of stellar interiors. About this time, electronic computers become more readily available, and Schwarzschild’s book spurred increased use of them to rapidly calculate many models for the interiors of stars of different masses and in different evolutionary stages. His basic methods can still be applied today to personal computers and even programmable calculators.

The year 1957 also saw the dawn of the Space Age with the launch of Sputnik 1, the first artificial Earth satellite, by the Soviet Union. Visible light is distorted as it passes through Earth’s turbulent atmosphere, resulting in images that are somewhat blurred. One way to get around this problem and produce sharper images is to observe from satellites above our atmosphere. The Hubble Space Telescope, placed in orbit at an Altitude of 600 kilometers in 1990, has supplied incredible high-resolution images of various stages of stellar evolution. Since observing time on space telescopes is limited and sparingly assigned, techniques have been developed for ground-based telescopes to provide sharper images. In speckle interferometry, many short exposure images (with little blurring due to atmospheric turbulence, since the exposure is so short) are recorded electronically and then combined in a computer to produce a sharper final image. In adaptive optics, sensors monitor the effects of atmospheric turbulence and control small motors that slightly alter the telescope optics to compensate for it and reduce the blur. These techniques have given new life to large ground-based telescopes.

Observations from space are essential, however, for observations in certain nonvisible portions of the electromagnetic spectrum, because the Earth’s atmosphere is opaque to much of the electromagnetic spectrum, including gamma rays, X-rays, and most ultraviolet and infrared radiation. Gamma rays, X-rays, and ultraviolet radiation are emitted by hot objects; Infrared radiation is emitted by cool objects. It is necessary to observe in all these wavelength regions to study both the cool births and the hot deaths of stars. Placing astronomical instruments on board satellites orbiting outside Earth’s atmosphere has allowed astronomers to observe over all parts of the electromagnetic spectrum. Infrared observations are revealing regions of star formation and the embryonic stars developing there. Around many of them it is possible to detect disks in which planets might eventually form. Ultraviolet and X-ray observations have revealed that stars lose much more mass during the giant and supergiant stages of their lives than was previously thought, so computer models of stellar evolution are being revised to account for this discovery.

Bibliography:

Asimov, Isaac. The Exploding Suns: The Secrets of the Supernovas. New York: E. P. Dutton, 1985.

Chaisson, Eric, and Steve McMillan. Astronomy Today. 6th ed. New York: Addison-Wesley, 2008.

Cohen, Martin. In Darkness Born: The Study of Star Formation. New York: Cambridge University Press, 1988.

Cooke, Donald A. The Life and Death of Stars. New York: Crown, 1985.

Dobrijevic, Daisy. "Solar System Planets, Order and Formation: A Guide."Space.com, 29 Mar. 2023, www.space.com/16080-solar-system-planets.html. Accessed 28 Sept. 2023.

Fraknoi, Andrew, David Morrison, and Sidney Wolff. Voyages Through the Universe. 3d ed. New York: Brooks/Cole, 2006.

Freedman, Roger A., and William J. Kaufmann III. Universe. 8th ed. New York: W. H. Freeman, 2008.

Genet, Russell M., Donald S. Hayes, Douglas S. Hall, and David R. Genet. Supernova 1987a: Astronomy’s Explosive Enigma. Mesa, Ariz.: Fairborn Press, 1988.

Greenstein, George. Frozen Star. New York: Charles Scribner’s Sons, 1983.

Kippenhahn, Rudolf. One Hundred Billion Suns: The Birth, Life, and Death of the Stars. New York: Basic Books, 1985.

Mehta, Jatan. "Solar System History 101."The Planetary Society, 14 Jan 2021, www.planetary.org/articles/solar-system-history-101. Accessed 28 Sept. 2023.

"Our Solar System." NASA Solar System Exploration, 28 Sept. 2023, solarsystem.nasa.gov/solar-system/our-solar-system/in-depth/#. Accessed 28 Sept. 2023.

Schneider, Stephen E., and Thomas T. Arny. Pathways to Astronomy. 2d ed. New York: McGraw-Hill, 2008.