Flowchart of Stellar Evolution

• All stars follow the same basic series of steps in the lives.
• Low-mass stars go through a red giant phase which ultimately turns into a planetary nebula with a white dwarf in the center.
• High-mass stars go through a red supergiant phase which ultimately results in a supernova, leaving behind either a neutron star or a black hole.
• The high-mass stars burn brighter and burn out faster, as depicted in this mass-luminosity relation.
• The deciding factor in the fate of a star is its mass.
  • Stars whose core is less than 1.4 solar masses (the Chandrasekhar limit) will leave behind a white dwarf, the size of which is inversely related to its mass. Note that the initial mass of this star is much greater than its core, but much of the mass is lost once the planetary nebula separates from the core.
  • Stars whose core is in the range 1.4-3 solar masses will leave behind a neutron star (much denser than a white dwarf). Note that the initial mass of this star is much greater than its core, but much of the mass is lost during the supernova phase.
  • Above 3 solar masses (the Tolman-Oppenheimer-Volkoff limit), a quark star might be created, although this is currently mostly conjecture.
  • Any stellar core over 5 solar masses will inevitably succumb to gravitational collapse, producing a black hole (much denser than a neutron star).
  • While a greater initial mass generally results in a greater remnant mass, there is some variability in the pathway. For example, midrange stars (between 5 Msun and 10 Msun) can end up either in a planetary nebula/white dwarf or a supernova/neutron star.

Stellar lifetimes

The lifetimes of star range from around a million years for the blue giants to billion of years for the red dwarfs.

Nova

• If a white dwarf has a close binary companion, the white dwarf may accrete gas from the companion's outer atmosphere. The gravitational energy released by the captured gas may be sufficiently great to start a fusion reaction on the surface of the white dwarf. This flare-up is known as a nova.
• A nova
  • can recur many times.
  • is about 100,000 solar luminosities.
  • fades after a few months (sometimes years).
  • is more common than supernovas (2-3/yr observed, about 200 so far).
• Eventually, the white dwarf could explode as a Type Ia supernova if it absorbs enough mass from its companion to push it over the Chandrasekhar limit.

X-ray burster (or accretion-powered pulsars)

• Like accreting white dwarfs that occasionally flare up into novae, an accreting neutron star (i.e., a neutron star which is a member of a binary system) can flare up with enormous luminosity.
• The infalling gas can reach half the speed of light before it impacts the neutron star surface. So much gravitational potential energy is released by the infalling gas, that the hotspots can have a peak luminosity nearly a hundred thousand times that of our Sun. This is the equivalent of detonating the entire world's nuclear arsenal on every square centimeter of the neutron star's surface within a minute.
• Temperatures of millions of degrees are produced, which results mostly in x-ray radiation, hence the label x-ray burster.
• Gas is accreted from the stellar companion is channeled by the neutron star's magnetic field on to the magnetic poles producing two or more localized X-ray hot spots similar to the two auroral zones on the Earth but far hotter. As the neutron star rotates, pulses of X-rays are observed as the hotspots move in and out of view if the magnetic axis is tilted with respect to the spin axis.
• The X-ray periods range from as little as a fraction of a second to as much as several minutes.

Supernova

• There are two very different types of supernova: Type I and Type II. These two types differ by their light curves and by their composition.
• Type II (also known as core-collapse) supernova:
  • A massive star evolves and eventually runs of fuel. The star then collapses and releases a lot of gravitational energy, which results in a massive supernova explosion.
  • The star's outer envelope expands and cools as it is blown into space by the shock wave sweeping up from below. This produces the characteristic "plateau" in the light curve for a few months after the maximum.
  • The expanding material consists mainly of unburned gas—hydrogen and helium—so it is not surprising that those elements are strongly represented in the supernova's observed spectrum.
  • Betelgeuse is an example of a Type II supernova.
    • If Betelgeuse (which is about 640 light-years away from us) went supernova, it would be over 10 times as bright as the full moon in our sky and that we would be able to see it even during the daytime for a few weeks.
• Type I (also known as thermal runaway) supernova:
  • Occurs in a binary system when one or both of the companions is a white dwarf.
  • When the white dwarf accretes enough mass from its companion to exceed the Chandrasekhar mass (1.4 solar masses), it collapses in a supernova explosion.
  • As this process unfolds, the internal temperature rapidly rises to the point at which the carbon-rich white dwarf can fuse carbon into heavier elements. This fusion provides the energy to detonate the white dwarf and produce the supernova.
  • In an alternative scenario, two white dwarfs in a binary system may collide and merge to form a massive, unstable star. The end result is the same—a carbon-detonation supernova.
  • Because the hydrogen-helium envelope has already been blown off in the planetary nebula, very little signal is detected from these two elements.
  • Because all Type I supernovae orginate from white dwarfs with the same (Chandrasekhar) mass, the energy released is very consistent and these are thus used as standard candles (to measure distance).

• Interesting facts:
  • Our solar system may have started with a nearby supernova.
  • Supernovae enrich the universe with heavy elements.
  • Radiation from nearby (with a couple of hundred light-years) supernovae helps to drive evolution.
  • Some supernovae can give rise to Gamma Ray Bursts (GRB's), which are so powerful that they could obliterate all life on Earth if they occurred in our vicinity.
    • This video tutorial describes GRB's in a user-friendly way.
  • This video tutorial summarizes the basic properties of supernovae.
• Frequency of occurrence:
  • A supernova occurs roughly once every 50 years in a galaxy roughly the size of the Milky Way.
  • With the aid of telescopes, more than 10,000 have been observed in other galaxies, but only a few in our own Milky Way Galaxy.
  • No supernova has been seen in the Milky Way since 1604.
  • The most reliable obsevations in the Milky Way are those of the years 1054, 1572, and 1604. Interestingly, not even one has been observed since the first use of the telescope for astronomy in 1609. The difficulty in seeing supernovas in our galaxy is that the dust blocks our view.
  • The last naked-eye supernova was observed in 1987, in a nearby galaxy (Large Magellanic Cloud).
• The best known supernova in the Milky Way occurred in ad 1054 and left behind a remnant known as the Crab Nebula.
• For a typical supernova to be as bright to the human eye as the Sun, it would have to occur at a distance of about one light-year from Earth. If Betelgeuse, which is about 500 light-years away, were to go supernova, it would rival the Moon in brightness.
• The nearest supernova witnessed in 4 centuries was Supernova 1987A, in the Large Magellanic Cloud (galaxy). This supernova was particularly useful to study because:
  • Its progenitor had been observed previously.
  • Its distance from us (168,000 ly) and its location in the Large Magellanic Cloud was already known.
  • It was the first opportunity for modern astronomers to see a supernova up close.
  • It was the closest observed supernova since SN 1604, which occurred in the Milky Way itself.
  • In fact, it was close enough to be seen with the naked eye (peak apparent magnitude 3) from the southern hemisphere.
  • It occurred after new telescopes, such as Hubble, could observe it in great detail.
  • It was the first opportunity for modern theories of supernova formation to be tested against observations.
  • It was the first time neutrinos emitted from a supernova had been observed directly.
• The Crab Nebula (M1) is the remnant of a supernova observed in A.D. 1054.
  • About 6000 light-years away, in the constellation Taurus.
  • Recorded as a "guest star" by Chinese astronomers.
  • It was visible with the naked eye for 23 days in daylight (about magnitude -6 or about 4 times brighter than Venus at its brightest) and 653 nights before fading from view.
  • Pulsar radiates energy from the center.
  • Rediscovered in 1758 by French astronomer Charles Messier (1730-1817), when he was looking for Comet Halley on its first predicted return. At first, Messier thought that it was a comet but recognized that it had no apparent proper motion. It was the discovery of this object which led Messier to compile his famous catalogue of nebulae and star clusters so that others would not confuse them with comets--beginning with the Crab Nebula as Messier object "Number 1" or M1.
• The undead star: The temperamental Eta Carinae brightened and dimmed over the last couple of centuries and may be ready to go supernova.
  • 8,000 light-years away
  • 100 times more massive than our Sun
  • 5 million times more luminous than our Sun
  • brighter than Rigel in 1837 and has been fluctuating wildly in brightness ever since
• The study of typical light curves from Type I supernovae indicates that radioactive nuclei form as a result of the explosion. These curves have two distinct features. After the initial peak, the luminosity first declines rapidly, then decreases at a slower rate. This change in the luminosity decay invariably occurs about 2 months after the explosion, regardless of the intensity of the outburst. Depending on the initial mass of the exploded star, the luminosity takes from several months to many years to decrease to its original value, but the shape of the decay curve is nearly the same for all exploded stars.
• We can explain the two-stage decline of the luminosity curve in terms of the radioactive decay of unstable nuclei, notably nickel-56 and cobalt-56, produced in abundance during the early moments of the supernova explosion. From theoretical models of the explosion we can calculate the amounts of these elements expected to form, and we know their half-lives from laboratory experiments. Because each radioactive decay produces a known amount of visible light, we can then determine how the light emitted by these unstable elements should vary in time. More direct evidence for the presence of these unstable nuclei was first obtained in the 1970s, when a gamma-ray spectral feature of decaying cobalt-56 was identified in a supernova observed in a distant galaxy.

Stellar nucleosynthesis

• The early Universe was hot and dense (like the core of a star) and nucleons could fuse together to make some helium. Aside from a tiny amount of lithium, nucleosynthesis beyond helium could not take because the Universe expanded and cooled too quickly. (The heavier the elements to be fused, the greater the required temperature.) Once all nuclear reactions ceased, 75% of the mass of the Universe was in the form of hydrogen and 25% in the form of helium. Today, after many cycles of stellar evolution, the Universe is still mostly hydrogen and helium, in roughly the same proportions, although 1% now fills out the rest of the periodic table. Your body is made mostly of these 'trace' elements.
• Stellar nucleosynthesis occurs at many different stages of stellar evolution, from main-sequence stars all the way to supernovae.
  • In perhaps the simplest nucleosynthesis reaction in the stellar core, helium is produced from hydrogen. This is accomplished by two known sets of fusion reactions: the proton-proton chain reaction and the CNO cycle.
  • In the proton-proton chain reaction, 4 hydrogen nuclei fuse to form helium. This reaction is followed by the fusion of 3 helium nuclei to form carbon-12. This conversion of three helium nuclei (alpha particles) to carbon-12 is called the triple-alpha process.
  • The element produced in the fusion of 2 helium nuclei is an unstable isotope of beryllium (Be-8). Once formed, most of the Be-8 rapidly decays back into two helium nuclei. A tiny fraction, however, survives long enough to capture another alpha particle to produce carbon-12, which is quite stable. (Because the triple-alpha process is unlikely, it requires a long period of time to produce carbon. One consequence of this is that no carbon was produced in the Big Bang because within minutes after the Big Bang, the temperature fell below that necessary for nuclear fusion.)
  •  The second of two known sets of fusion reactions by which helium is produced from hydrogen is the CNO cycle. In this cycle, there is still a net production of helium from hydrogen, but carbon, nitrogen, and oxygen isotopes act as catalysts (i.e., are regenerated in each cycle) to facilitate the reaction.
  • In stars less massive than the Sun, the proton-proton chain is more dominant. In stars more massive than the Sun, the CNO cycle is dominant. In general, the higher the temperature, the more dominant the CNO cycle is relative to the simple proton-proton chain reaction.
  • Elements heavier than helium but lighter than iron are built predominantly through helium capture. These elements have masses which are divisible by 4, the mass of a helium nucleus, and are found in the peaks of the chart of cosmic abundances.
  • Although many different nuclear reactions undoubtedly occur in the cores of evolving stars, the efficiency of fusion decreases with increasing size of the reacting nuclei because of increased internuclear repulsion. For example, the repulsive force between two carbon nuclei is three times greater than the repulsive force between a nucleus of carbon and one of helium. Thus, at a given temperature, carbon-helium fusion should occur more efficiently than carbon-carbon fusion. Heavier elements therefore tend to form predominantly through helium-capture rather than through fusion of heavier-than-helium nuclei, though these reactions occur as well. Consistent with theory and observation, elements with atomic masses divisible by 4 (i.e., which are built up through helium-capture)--oxygen-16, neon-20, magnesium-24, silicon-28, etc.--stand out as prominent peaks in the chart of cosmic abundances. Helium, of course, stands out as well since it was produced in abundance in the Big Bang and continues to be produced in the cores of evolving stars.
  • Other reactions for pre-iron nucleosynthesis are possible but are also less common than helium capture. For example, protons and neutrons are freed from some nuclei and are absorbed by others, resulting in new nuclei with masses intermediate between those formed by helium capture. Such intermediate elements (e.g., fluorine-19, sodium-23, phosphorus-31, and many others) have masses which are not divisible by 4, the mass of a helium nucleus, and are found in the troughs of the chart of cosmic abundances.
  • Around the time silicon-28 appears in the core of a star, the star's core temperature has reached an incredible 3 billion K, and the gamma rays associated with this temperature have enough energy to break a nucleus apart (photodisintegration). This leads to a competitive struggle between the continued capture of helium to produce even heavier nuclei and the tendency of more complex nuclei to break down into simpler ones. This is the same process of photodisintegration will ultimately accelerate the collapse of the star's iron core toward a Type II supernova.
  • The process of photodisintegration provides raw material that allows the helium-capture process to proceed to greater masses. The process continues, with some heavy nuclei being destroyed and others increasing in mass. In succession the star forms sulfur-32, argon-36, calcium-40, titanium-44, chromium-48, iron-52, and nickel-56. This two-step process—photodisintegration followed by the direct capture of some or all of the resulting helium-4 nuclei (or alpha particles)—is often called the alpha process.
  • Because nickel-56 is unstable, it decays rapidly, first into cobalt-56, then into a stable iron-56 nucleus. Any unstable nucleus will continue to decay until stability is achieved, and iron-56 is the most stable of all nuclei. Thus, the alpha process leads inevitably to the buildup of iron in the stellar core.
  • Elements heavier than iron cannot be produced through helium capture because they are very unstable and decay back to iron, the most stable element, as quickly as they are formed. Iron's 26 protons and 30 neutrons are bound together more strongly than the particles in any other nucleus. Iron is said to have the greatest nuclear binding energy per nucleon of any element and is thus the most stable element. This enhanced stability of iron explains why nuclei tend to "accumulate" near iron as stars evolve.
  • Elements beyond iron form by neutron capture and radioactive decay. The heaviest nuclei of all are formed by an extremely intense, but brief, period of neutron capture during a type II (core-collapse) supernova explosion. Because the time available for synthesizing these heaviest nuclei is so brief, they never become very abundant. Consequently, elements heavier than iron are collectively about a billion times less abundant than hydrogen and helium.
• Two types of neutron capture are recognized: the s-process and the r-process.
  • The s-process or slow-neutron-capture-process refers to a series of nuclear reactions which occur at relatively low neutron density. The s-process produces approximately half of the isotopes of the elements heavier than iron, up to the element Bismuth-209, which is the heaviest-known nonradioactive nucleus. Nuclei more massive than that of Bismuth-209 cannot be built up through neutron capture in the low-neutron-density environment of the s-process; these nuclei are too unstable and decay back to Bismuth as quickly as they are formed. The s-process is responsible for the synthesis of copper, lead, silver, and gold, among others, as well as some of the lighter elements intermediate between those formed by helium capture.
  • In the r-process or rapid-neutron-capture-process refers to a series of nuclear reactions which occur in the high-neutron-density environment of a Type II supernova. The neutron-capture rate in this environment is so great that even the most unstable heavy nuclei can capture many neutrons before decaying. Virtually all the super-heavy elements in the Universe, including radium, thorium, uranium and plutonium, are produced via the r-process.
• Evidence for nucleosynthesis in the stars includes:
  • Observed abundances of various elements and their isotopes in the cosmos are in strong agreement with theoretical predictions based on known properties of nuclear reactions.
  • The observed presence of relatively short-lived technetium-99 in the spectra of many red-giant stars was one of the first strong pieces of evidence that heavy elements really do form in the cores of dying stars. Because technetium is radioactive, with a half-life (0.2 million years) much less than the age of the star, its abundance must reflect its creation within that star during its lifetime.
  • The shape of light curves of supernovae is consistent with theoretical expectations.

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