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Continuing with our study of the lives of stars, we follow them through their adolescence, old age, and their eventual death. Depending upon the mass they had at birth, some stars live an extremely long time while others go through life fast and furious. The stars that live a long time die without much fanfare compared to the stars that live only a relatively short period of time. The massive, short-lived stars go out in a blast of energy representing one of the most violent explosions in the Universe. We've addressed the question, "How do we know all this?" by having the student "walk" through the process of constructing what is known as a "color-magnitude diagram" of two clusters of stars. These diagrams reveal not only the ages of the clusters (and thus the way stars of different masses age) but also the relative distances to the clusters.
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After completing this lesson, you should be able to
You are at the point in your study of astronomy where you should be asking yourself: "How do we know all of this?" We cannot go to the center of the Sun, we cannot travel to the stars and grab a bit of plasma. We have been studying astronomy with large telescopes and at multiple wavelengths for only a few decades, using high-powered computers to model stars for even less time. Even the lives of the shortest-lived stars are a million times longer than any human life. As the text mentions, "we cannot hope to see the life story of any single star unfold before our eyes or telescopes."
When working with the classification of stars on the H-R Diagram, we assumed that if a star had nearly the same spectrum as another star, these stars were, in fact, nearly identical. When working with how stars live and die, we assume that if a star has nearly the same mass as another star, then those two stars will follow similar evolutionary paths. Well, we still cannot follow any given star for millions or billions of years. What we can do, however, is study clusters of stars—stars that we assume are of nearly the same age and chemical composition (having formed from the same giant molecular cloud). With these two variables taken care of, the differences we see in temperatures and luminosities for stars in a cluster must be due to differences in their masses.
Before going on with these notes, read carefully through Chapters 13 and 14. Take a sheet of paper and jot down a few notes as you read. Note if the star is in equilibrium (gravity and gas pressure are balanced, see Section 7.3 for a review), or if either gravity or gas pressure is getting an "upper hand." When we get to stars that have finished fusing hydrogen in their cores and thus are "leaving" or "evolving off of" the main sequence, we need to look at different regions of the interior and not just the whole star. As a star evolves, its core starts evolving separately from the rest of the star, and the only thing the rest of the star can do is respond. The response of the rest of the star to what is happening in its core and just outside the core is truly amazing! You will find as you read through these two chapters, that stars like our Sun (we'll call them "low-mass stars") follow a different path from stars that are approximately six times or more the mass of our Sun (we'll call these "high-mass stars"). Keep track of these differences. How a star finally ends its life depends on its mass. A star can go relatively quiet, fading away for eternity, or it can go out with one of the most powerful explosions known in the Universe.
After you have finished reading the chapters for this lesson, and have worked through the review and thought questions at the end of each chapter (see the Self-Review section below for hints on these questions), then return to these notes and try to answer the questions at each stage of the flow chart based just on your reading and your notes. (You may also need to do a quick review of the birth of stars from Chapter 12, comparing low-mass and high-mass stars once more.) Then, "touch a star" and check your answer.
We start this lesson with a flowchart of stellar evolution. Since the mass of a star determines its entire life (as long as it does not belong to a close binary system where other factors come in to play), we can follow the various paths taken by stars of different masses. This flowchart comprises the first half of these notes.
So much of what we will be discussing here and that is covered in the text relies upon theory. We simply cannot look deep into any star and we certainly can't wait around for millions or billions of years to see how any given star ages and dies. The theory is supported, indeed driven by our observations of open and globular clusters. The observation of clusters completes these notes.
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Stars are born within molecular clouds in the Galaxy. As the cloud collapses it fragments, and multiple stars are formed. We see an open (galactic) cluster as a result. As the star collapses, its overall density increases. As the density increases, the temperature and pressure at the center of star increases as well, quite dramatically. At this point in the star's life, its luminosity is provided by release of gravitational energy. Low-mass stars may take about a million years in this collapse until they finally arrive on the main sequence. Massive stars are born from huge molecular clouds. As protostars, they collapse much more rapidly than their low-mass counterparts, maybe going through the whole process in 10,000 years. The profuse amount of radiation produced by these O and B stars drive the dynamics and subsequent star formation of the whole cloud. (Images: NASA/STSCI/AURA) |
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The mass of a protostar determines its place on the H-R diagram, its energy source, its ultimate fate. Low mass stars sit on the main sequence at low luminosities, low temperatures; high mass stars, at high luminosities and high temperatures. (Applying this to humans: it would be as if 6-pound babies would always grow up to be lawyers; 8-pound babies, doctors; 10-pound babies, Paris models; etc.) For purposes of this discussion, we will call stars with less than 3 times the mass of the Sun "low-mass" stars. We will call stars with masses greater than 6 times the Sun's "high-mass stars." Those stars with masses in the range of 3-6 times the Sun don't get paid much attention here. Although they follow much of the story of a high-mass star, they do not get to have the catastrophic explosion at the end. As you will see, though, they do become somewhat rare objects.
A Brown DwarfIf the star does not have enough mass to create high enough temperatures and pressures in its core, then it can never commence hydrogen fusion. It may spend a brief time fusing deuterium (heavy hydrogen), but ultimately its only source of luminosity is the heat released during its gravitational collapse. It is doomed to cool for a very long time and eventually turn into a cold mass. If Jupiter had about 80 times its current mass, we would have a glowing, brown dwarf in the sky. (Image: NASA/STSCI/AURA)
As we learned when we studied the Sun in Lesson Five, energy generation takes place in the core via the proton-proton cycle: fusing 4 hydrogen nuclei (protons) into 1 helium nucleus (2 protons, 2 neutrons), releasing energy in the process. The famous equation that explains where energy comes from, E = mc2, tells the whole story of how a little bit of mass can produce a whole lot of energy. Stars' lifetimes on the main sequence depend on their masses. Low mass stars spend 10 billions years or more, while high mass stars may stay for a mere 3-4 million years. The star's luminosity while on the main sequence is described by the mass luminosity relationship: L is proportional to M3.2. What this means is that even though high-mass stars have much more fuel than low-mass stars, they go through it at a much, much greater rate. (An analogy here on Earth: getting 70 mpg and having a 10-gallon tank versus getting 10 mpg and having a 20-gallon tank. The efficient car can go 700 miles on a tank of gas, while the low-mileage car barely covers 200 miles.) While on the main sequence, the star is in hydrostatic equilibrium: the outward pressure exerted by hot plasma—heated by the energetic photons produced in the core—balances the inward pressure due to the force of gravity. For stars roughly the mass of our star, outside the core lies the radiative envelope where the photons partake in a "random walk". It takes roughly a million years for a photon to get from the core (as a gamma ray) to the surface (as a visible light photon). The photon travels a mere 1 cm before being scattered by an electron or an ion. Outside the radiative envelope lies the convection zone (convective envelope). Heat is transferred here by the bulk movement of material. The top of the convection zone displays granulation at the base of the atmosphere. The atmosphere (if the star is similar to our Sun) has a photosphere, a chromosphere, and a corona. Many stars have huge starspots on their surface that can be detected by telescopes here on Earth.
This
field of stars in the constellation of Sagittarius shows stars of many colors,
including those that are orange and red. These stars are the "red giant
stars." (Image
courtesy Hubble
Heritage Team and NASA)
The cores of red giant stars are composed primarily of helium, ashes from the billions of years of fusion of hydrogen. The temperature in the core is not yet high enough to fuse elements heavier than helium. Without any support from energy generation, the helium core contracts. Here is where we can bring in our basic knowledge of physics: the material in the outer core has gravitational potential energy (GPE). When the core contracts, this GPE releases a profusion of energy—kinetic and thermal. The heat from this contraction causes the hydrogen in a shell surrounding core to ignite and we get "hydrogen-shell fusion." The core will continue to contract and release energy until it cannot contract any further without squeezing the electrons together. The core is now supported by electron degeneracy—no two electrons can exist in the same state, and resist doing so with enough force to halt the collapse.
The fusion of hydrogen to helium in the shell and the radiation pressure that results supports the rest of the star from gravitational collapse. The hydrogen shell produces a lot of energy. In a simplified interpretation of why stars turn into red giants (astronomers aren't really sure why this happens), we assume that because the hydrogen shell is closer to the surface of the star and the luminosity of that shell is very high, the outer layers of the star expand. The star becomes a red giant.
The star itself does not gain or lose mass (at least it has not lost much mass up to this stage) during its whole life, unless it is in a close-binary system where the stars are so close they are influencing each other dramatically. But, as the hydrogen-fusing shell works its way out through the interior of the star, it continuously "rains" helium ash onto the core of the star. The core is slowly gaining mass at the expense of the rest of the star. This extra mass is squeezing the core, raising the temperature there and steadily trying to raise the density.
Helium Flash! and Horizontal BranchAt left is a Hubble Space Telescope of the globular cluster 47 Tucanae. The large number of red giant stars in this cluster is noticeable by the distinctly yellowish-orange color of these stars. (Image courtesy W. Keel and NASA)
Helium core fusion starts in the core of a red giant when the core temperature and density reaches 108 K and 108 kg/m3. As a red giant, the core was supported by electron degeneracy. Degeneracy is a strange state of matter indeed. Normal rules of physics do not apply here. A degenerate stellar core does not follow normal gas laws that dictate if something shrinks, the pressure and temperature rise, or if something expands, its pressure and temperature decreases. Once there is enough mass dumped on the core from the rest of the star, fusion starts in a runaway explosion of the entire core simultaneously. The red giant experiences a "helium flash." Unfortunately for us, we cannot see this flash; in fact, the rest of the star doesn't even know it happened until there appears to be a stable source of radiative support again. The degeneracy in the core disappears since fusion is now taking place. The star is supported by the energy produced via helium-to-carbon fusion in the core. The luminosity of the star decreases; the temperature increases. The star is in a second zone of stability, the horizontal branch, that will last a few billion years.
Not all stars are totally stable in this region of the H-R Diagram. There are exceptions, stars that pulsate. These stars are called RR Lyrae Stars, and we will be getting better acquainted with these variable stars and what they have revealed about our galaxy in the lab that accompanies Lesson Eleven.
Eventually,
over a time frame shorter than the hydrogen-fusion stage, the core of the horizontal-branch
star will eventually run out of helium to fuse into carbon. Fusion shuts down,
and the core starts to contract, having no means of supporting itself. This
small, contracting core is made of carbon ash. The gravitational contraction
releases the energy that had been stored as potential energy, igniting a helium-fusing
shell around it. Luminosity continues to be generated by the shrinking core
until it again cannot contract anymore due to the onset of electron degeneracy.
Meanwhile, back to the rest of the star: the rest of the star expands again
and becomes an asymptotic giant branch star, almost retracing its path as a
red giant star (see the link to the H-R diagram for the globular cluster 47
Tucanae shown to the right). The helium fusing shell surrounding the core rains
down "carbon" ash. The luminosity produced by the helium-to-carbon
fusion in this shell ignites a hydrogen-to-helium fusing shell on the outside
of it. This is called a double-shell-burning star. The hydrogen shell is raining
down helium ash to the helium-burning shell, while the helium shell is raining
down carbon ash to the core. The two H and He fusing shells support the outer
parts of the star. Outer atmosphere expands even more than before.
Planetary NebulaEnergy production in the star will eventually stop as the He- and H-fusing shells work their way through the star. At this stage, the upper envelopes and the atmosphere are extremely unstable. Large pulsations in the star have driven almost half of the star's material into interstellar space, leaving behind only the once-active stellar core. This carbon core, supported now strictly by electron degeneracy, is extremely dense: 1010 kg/m3. The gaseous object surrounding the core is called a planetary nebula. It glows because of the energetic, ionizing photons produced by the now-exposed, extremely hot core. (Planetary Nebula Mz 3: Image courtesy Hubble Heritage Team and NASA)
The star has finally reached the twilight of its years and has become a white dwarf: an immensely hot, Earth-sized object. What's left of a once magnificent star is doomed to a slow, slow death. What we see is the remnant of the core of the star. Its fate is to cool for billions and billions of years until it becomes a black, ultra-cold cinder. The carbon and other elements it contains will never be a part of the life of a star again.
Massive white dwarfs are produced by stars whose mass lies between roughly three and six times the mass of the Sun. These white dwarfs contain mostly oxygen and neon, as well as some carbon, representing the fact that the star they came from was able to fuse elements slightly heavier than helium. These white dwarfs are rare as their mass lies right on the limit for electron degeneracy to be able to support the object from collapsing. If the white dwarf has a mass of slightly more than 1.4 times that of the Sun (the Chandrasekhar limit), it collapses into a neutron star.
NGC 7635, the Bubble Nebula, is 10 light-years across, more than twice the distance from the Earth to the nearest star. The nebula is made up of an expanding shell of glowing gas surrounding a hot, massive star in our Milky Way galaxy. This shell is being shaped by strong stellar winds of materials and radiation produced by the bright star at the left, which is 10 to 20 times more massive than our Sun. These fierce winds are sculpting the surrounding material, composed of gas and dust, into the curve-shaped bubble. (NGC 7635, the Bubble Nebula: Image courtesy Hubble Heritage and NASA)
High Mass Stars: On the Main SequenceMassive stars start their lives much like their solar counterparts, only they tend to over-do everything! They truly live fast and furious lives. In addition to the proton-proton cycle, massive stars can also produce energy by converting hydrogen to helium via the C-N-O cycle. Interior pressures and temperatures are so great that helium and carbon are often fused along with hydrogen, fusion isn't limited to just one process. One of the most massive stars ever observed, being about 150 times more massive than the Sun (or, theoretically, about as massive as any star can ever be), Eta Carinae is prone to massive outbursts. The most massive stars, like Eta Carinae, produce so much luminosity, that they literally blow much of their outer layers away early on in life. These stars have episodes where they brighten by many magnitudes and spew material into interstellar space. We expect this particular star, which may in fact be a binary system, to really "blow its top" some day. (Image courtesy of J. Hester and NASA)
Massive stars spend very little time on the main sequence. For the most massive stars, this is only a few million years. These stars move to luminous red supergiant region as hydrogen is exhausted in core. The core contracts slightly, but since the star already has high enough temperatures to fuse helium to carbon via the triple-alpha process (review Section 13.4 of the text), it moves to this stage relatively smoothly. Along with a helium-burning core, a hydrogen-to-helium fusing shell just outside of the core may also be present.
Helium Depletion in the CoreThe star has been primarily supported by helium-to-carbon fusion, but the helium has run out. The core again constracts. Some of the carbon and oxygen get fused into neon. As each new fusion process cycles to another, the star changes its temperature and moves from a red supergiant to a blue supergiant, and perhaps back to red again. At this stage in the star's life it is basically in hydrostatic equilibrium, although fusion in the core becomes dramatically dependent upon the temperature there. Any slight changes in the temperature results in extreme changes in the fusion rate. The carbon-to-oxygen-to-neon fusing core is surrounded by helium-to-carbon fusing shell which is surrounded by a hydrogen-to-helium fusing shell. As the fusion rates increase and decrease, the star's atmosphere responds. It is not unusual to find a red supergiant that changes its luminosity by 100 or more times (5 magnitudes) at this stage in its life. The image at the left is of the red supergiant, Betelgeuse, in Orion. (Image courtesy of A. Dupree and NASA)
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For these stars having masses originally greater than six times the Sun's, their outrageous "partying" isn't over yet. High-mass stars proceed to fuse oxygen to neon, neon to magnesium, magnesium to silicon, and finally silicon to iron (with intermediate steps along the way). Each stage leaves a "shell" of the previous stage fusing above it, raining ash onto the present core. Each stage happens progressively faster than the previous ones. The star is now caught in the cycle of depletion of one element to fuse in its core, contraction of the core, heating, fusion of the next heavier element, igniting of another fusing shell. Star "loops" across high-luminosity part of the H-R Diagram as each stage progresses.
Figure 13.4 of your text does not show the looping of a high mass star. To the right, is a modified version of that figure that shows approximately what happens.
When does fusion stop? The element iron has a particularly unique characteristic: its nucleus is the most tightly bound of any atom. Energy has been released as fusion as progressed from hydrogen to iron, but once an iron core is produced, fusion stops. No energy is forthcoming from fusing iron into a heavier element. Beyond iron, energy can be obtained only by fission. The star suddenly finds itself without any means of support against the force of gravity at all!
With no further support being provided for the star, gravity must win. The
star starts to implode. The energy stored as gravitational potential energy
is released suddenly. This energy results in the nuclei of the elements being
split apart in a process called photodisintegration. All the millions of years
the star spent in fusing heavier and heavier elements was for naught. Even more
energy disappears in this process—since energy was produced during the
fusion, energy is needed for the splitting. The collapse accelerates as there
is even less energy available for support. It gets worse! Protons and
electrons are crushed together under extreme, unbelievable really, temperatures
and pressures and form neutrons. Neutrinos are also formed in this merging,
carrying away even more energy directly from the star. Electron degeneracy simply
cannot act to support the core. 
The
density of the core has now reached 1015 kg/m3. Just as a rubber
ball will compress before it bounces back, the stellar core overshoots and rises
to an astounding 1018 kg/m3.(There just are not enough superlatives
in the English language to describe what is happening inside the dying star!)
The core now rebounds with a vengeance, ricocheting shock waves and material
back into interstellar space. This whole process takes place in less time than
it took you to read this paragraph.
Supernova 1987a provided observational support to the theory of these explosions. The Earth received a pounding of neutrinos first, although only a small fraction of the billions and billions passing through the Earth were detected. Since neutrinos do not interact with matter, they were able to escape the collapse first. The light from the explosion followed shortly afterward. (Supernova 1987a: Image courtesy Hubble Heritage and NASA)
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The Crab Nebula is an example of what happens to the material from the supernova explosion. This supernova explosion happened in 1054 AD; the material is still speeding away from the neutron star left behind at 1000's of kilometers per second. Massive stars give almost all of their material back to interstellar space. Out of the 6-plus solar masses the star originally was born with, only between 1.4 and 3 solar masses remain in the neutron star (we deal with stellar black holes later). The neutron star is supported by "neutron degeneracy," a state similar to electron degeneracy, only much, much more dense. What remains of the star is essentially one huge atomic nucleus made of neutrons alone, an object about the size of the Earth but with an unbelievably high density. We will be examining the neutron star in the Crab Nebula more closely in Lesson Ten.
As the star explodes, there is also an abundance of neutrons spewing away from the neutron star. These neutrons run into nuclei of atoms and build up. When a neutron decays into a proton, a heavier element results. This process occurs very rapidly, with neutrons building up faster than they decay. In this way, the heaviest elements found on the periodic chart are formed—we can think of no other way for them to be produced. (Think about the definition of alchemy—if humans could turn baser elements into gold, we would have done it by now.) In addition to this rapid neutron capture, the shock waves proceed through the expanding stellar material bringing densities, pressures, and temperatures high enough to fuse elements, and to force helium nuclei into other nuclei, building up the periodic chart even faster. Thus, the formation of the elements found in the Universe is completed.
All pulsars are neutron stars, but not all neutron stars are pulsars. Observationally, a pulsar is an object that emits flashes of light several times per second or more, with near perfect regularity. Theoretically, pulsars are rapidly rotating neutron stars. What's the strongest evidence we have that pulsars are neutron stars? No massive object, other than a neutron star, could spin as fast as we observe pulsars spin.
The gold you wear, the nickel in our coins, the uranium in the ground, silver, mercury, iron, copper, platinum—all were once atoms in a stupendous explosion. Carbon, nitrogen, oxygen, silicon, and the other essential chemicals for life were also produced in a supernova. Thus are we truly the children of the stars.
Blackholes and their environment are so fascinating that they get a chapter all to themselves. You will have to wait until Lesson Ten to learn more about these intriguing and enigmatic objects.
From the textbook:
Learn more about the evolution of stars from a number of excellent sites:
