Astronomy – The Stars
Lifetime of a Star

A low-mass star (less than about 8% of the mass of the Sun) cannot produce helium from hydrogen though it may be able to fuse deuterium. It becomes a brown dwarf or, if less than about 13 times Jupiter’s mass, a sub-brown dwarf (if it orbits a larger star it is a planet). A brown dwarf is not necessarily brown, but more likely to be magenta or orange. Gas-giant planets have some of the characteristics of brown dwarfs. For example, Jupiter and Saturn are both made primarily of hydrogen and helium, like the Sun. Saturn is nearly as large as Jupiter, despite having only 30% the mass. Three of the giant planets in the Solar System (Jupiter, Saturn, and Neptune) emit much more heat than they receive from the Sun. And all four giant planets (these and Uranus) have their own “planetary systems” – their moons. Brown dwarfs form independently, like other stars, but lack sufficient mass to “ignite”. Like all stars, they can occur singly or in close proximity to other stars. Some orbit stars and can, like planets, have eccentric orbits.

For a protostar more massive than a brown dwarf, the core temperature will eventually reach 10 million K, initiating the proton-proton chain reaction and allowing hydrogen to fuse, first to deuterium and then to helium. In stars of slightly over one solar mass, the carbon–nitrogen–oxygen fusion reaction (CNO cycle) contributes a large portion of the energy generation. The onset of nuclear fusion leads relatively quickly to a hydrostatic equilibrium in which energy released by the core exerts a “radiation pressure” balancing the weight of the star’s matter, preventing further gravitational collapse. The star thus evolves rapidly to a stable state, beginning the main-sequence phase of its evolution.

A mid-sized yellow dwarf star, like the Sun, will remain on the main sequence for about 10 billion years. The Sun is thought to be in the middle of its lifespan; it is currently on the main sequence.

Nuclear fusion powers a star for most of its life. Initially the energy is generated by the fusion of hydrogen atoms at the core of the main-sequence star. Later, as the preponderance of atoms at the core becomes helium, stars like the Sun begin to fuse hydrogen along a spherical shell surrounding the core. This process causes the star to gradually grow in size, passing through the subgiant stage until it reaches the red giant phase. Stars with at least half the mass of the Sun can also begin to generate energy through the fusion of helium at their core, whereas more massive stars can fuse heavier elements along a series of concentric shells.

Once a star like the Sun has exhausted its nuclear fuel, its core collapses into a dense white dwarf and the outer layers are expelled as a planetary nebula (which has absolutely nothing to do with planets). Stars with around ten or more times the mass of the Sun can explode in a supernova as their inert iron cores collapse into an extremely dense neutron star or black hole. Although the universe is not old enough for any of the smallest red dwarfs to have reached the end of their lives, stellar models suggest they will slowly become brighter and hotter before running out of hydrogen fuel and becoming low-mass white dwarfs.

Eventually the core exhausts its supply of hydrogen and the star begins to evolve off of the main sequence of the Hertzsprung–Russell Diagram. Without the outward pressure generated by the fusion of hydrogen to counteract the force of gravity the core contracts until either electron degeneracy becomes sufficient to oppose gravity or the core becomes hot enough (around 100 MK) for helium fusion to begin. Which of these happens first depends upon the star’s mass.

Low-mass stars may stay on the main sequence for some six to twelve trillion years, gradually increasing in both temperature and luminosity, and take several hundred billion more to slowly collapse into a white dwarf. Such stars are fully convective and will not develop a degenerate helium core with hydrogen burning shells, or at least not until almost the whole star is helium, so they never expand into a red giant.

Slightly more massive stars do expand into red giants, but their helium cores are not massive enough to reach the temperatures required for helium fusion. When hydrogen shell burning finishes, these stars become white dwarfs.

A star like the Sun will become a red giant before expanding further to become a bright red giant with a luminosity thousands of times that of the current Sun. Its interior structure is characterized by a central and inert core of carbon and oxygen, a shell where helium is undergoing fusion to form carbon (known as “helium burning”), another shell where hydrogen is undergoing fusion forming helium (known as “hydrogen burning”) and a very large envelope of material of composition similar to main-sequence stars.

Mid-sized stars of roughly 0.5 to 10 solar masses become red giants, which are red and of a high luminosity. Examples include Aldebaran in the constellation Taurus and Arcturus in the constellation Boötes. Red giants all have inert cores with hydrogen-burning shells, concentric layers still fusing hydrogen into helium; their cores are made of helium or carbon. The accelerated fusion in the hydrogen-containing layer immediately over the core causes the star to expand. This lifts the outer layers away from the core, reducing the gravitational pull on them, and they expand faster than the energy production increases. This causes the outer layers of the star to cool, which causes the star to become redder.

The core’s gravity compresses the hydrogen in the layer immediately above it, causing it to fuse faster than normally. This in turn causes the star to become more luminous (from 1,000 to 10,000 times brighter) and expand; the degree of expansion outstrips the increase in luminosity, causing the effective temperature to decrease. The expanding outer layers of the star are convective, with the material being mixed by turbulence from near the fusing regions up to the surface of the star. For all but the lowest-mass stars, the fused material has remained deep in the stellar interior until this point, so the convecting envelope makes fusion products visible at the surface for the first time. At this stage of evolution, the results are subtle, with the largest effects, alterations to the isotopes of hydrogen and helium, being unobservable. The effects of the CNO cycle appear at the surface, with lower ¹²C/¹³C ratios and altered proportions of carbon and nitrogen. These are detectable with spectroscopy and have been measured for many evolved stars.

After a star has consumed the helium at its core, fusion continues in a shell around a hot core of carbon and oxygen. The star follows burns with even faster energy generation (which lasts for a shorter time). Although helium is being burnt in a shell, the majority of the energy is produced by hydrogen burning in a shell closer to the surface of the star. Helium from these hydrogen burning shells drops towards the centre of the star and periodically the energy output from the helium shell increases dramatically. Depending on mass and composition, there may be several to hundreds of such thermal pulses.

There may be a deep convective zone that brings carbon from the core to the surface, and can occur several times. In this way a carbon star is formed, very cool and strongly reddened stars with strong carbon lines in their spectra. A process known as hot bottom burning may convert carbon into oxygen and nitrogen before it can be dredged to the surface, and the interaction between these processes determines the observed luminosities and spectra of carbon stars in particular clusters.

Some stars, the Mira variables, pulsate with well-defined periods of tens to hundreds of days and large amplitudes up to about 10 magnitudes (in the visual part of the spectrum, total luminosity changes by a much smaller amount). In more massive stars the stars become more luminous and the pulsation period is longer, leading to enhanced mass loss, and the stars become heavily obscured at visual wavelengths. These stars can be observed as OH/IR stars, pulsating in the infra-red and showing OH maser activity. These stars are clearly oxygen rich, in contrast to the carbon stars, but both must be produced by dredge ups. [An OH/IR star has strong hydroxyl (OH) masers (“microwave amplification by stimulated emission of radiation”) and strong infrared (IR) emission from the shell of warm gas.]

These mid-range stars ultimately run out of fuel for shell burning. They are not sufficiently massive to start full-scale carbon fusion, so they contract again, going through a period that produces a planetary nebula with an extremely hot central star which then cools to a white dwarf. The expelled gas is relatively rich in heavy elements created within the star and may be particularly oxygen or carbon enriched, depending on the type of the star. The gas builds up in an expanding shell called a circumstellar envelope and cools as it moves away from the star, allowing dust particles and molecules to form. With the high infrared energy input from the central star, ideal conditions are formed in these circumstellar envelopes for maser excitation.

In Massive stars, the core is already large enough at the onset of the hydrogen burning shell that helium ignition occurs at an early phase. So when these stars expand and cool, they do not brighten as much as lower-mass stars, but as they were much brighter than lower-mass stars to begin with, they are still brighter than red giants formed from less massive stars. These stars are unlikely to survive as red supergiants; instead they will destroy themselves as type II supernovae.

Extremely massive stars (more than approximately 40 solar masses) are very luminous and have very rapid stellar winds; they lose mass so rapidly that they may strip off their own envelopes before they can expand to become red supergiants, and thus retain extremely high surface temperatures (and blue-white color). The largest stars of the current generation are about 100 to 150 solar masses because the outer layers would be expelled by the extreme radiation. Although lower-mass stars normally do not burn off their outer layers so rapidly, they can likewise avoid becoming red giants or red supergiants if they are in binary systems close enough so that the companion star strips off the envelope as it expands, or if they rotate rapidly enough so that convection extends all the way from the core to the surface, resulting in the absence of a separate core and envelope due to thorough mixing. The core grows hotter and denser as it gains material from fusion of hydrogen at the base of the envelope. In all massive stars, electron degeneracy pressure is insufficient to halt collapse by itself, so as each major element is consumed in the centre, progressively heavier elements ignite, temporarily halting collapse. If the core of the star is not too massive (less than approximately 1.4 Suns, taking into account mass loss that has occurred by this time), it may then form a white dwarf (possibly surrounded by a planetary nebula) composed chiefly of oxygen, neon, and magnesium.

Above a certain mass (estimated at approximately 2.5 solar masses and whose star’s progenitor was around 10 solar masses), the core will reach the temperature (approximately 1.1 gigakelvins – that’s 10⁹ degrees) at which neon partially breaks down to form oxygen and helium, the latter of which immediately fuses with some of the remaining neon to form magnesium; then oxygen fuses to form sulphur, silicon, and smaller amounts of other elements. Finally, the temperature gets high enough that any nucleus can be partially broken down, most commonly releasing an alpha particle (helium nucleus) which immediately fuses with another nucleus, so that several nuclei are effectively rearranged into a smaller number of heavier nuclei, with net release of energy because the addition of fragments to nuclei exceeds the energy required to break them off the parent nuclei.

A star with a core mass too great to form a white dwarf but insufficient to achieve sustained conversion of neon to oxygen and magnesium, will undergo core collapse (due to electron capture) before achieving fusion of the heavier elements. Both heating and cooling caused by electron capture onto minor constituent elements (such as aluminium and sodium) before collapse may have a significant impact on total energy generation within the star shortly before then. This may produce a noticeable effect on the abundance of elements and isotopes ejected in the subsequent supernova.

Supernovae: Once the nucleosynthesis process arrives at iron-56, the continuation of this process consumes energy (the addition of fragments to nuclei releases less energy than required to break them off the parent nuclei). If the mass of the core exceeds the Chandrasekhar limit, electron degeneracy pressure cannot support its weight against the force of gravity, and the core undergoes a sudden, catastrophic collapse to form a neutron star or (in the case of cores that exceed the Tolman–Oppenheimer–Volkoff limit) a black hole. Through a process that is not completely understood, some of the gravitational potential energy released by this core collapse is converted into a Type Ib, Type Ic, or Type II supernova. [See the article on supernovae for an explanation of the various types.] It is known that the core collapse produces a massive surge of neutrinos, as observed with supernova SN 1987A. The extremely energetic neutrinos fragment some nuclei; some of their energy is consumed in releasing nucleons, including neutrons, and some of their energy is transformed into heat and kinetic energy, thus augmenting the shock wave started by rebound of some of the infalling material from the collapse of the core. Electron capture in very dense parts of the infalling matter may produce additional neutrons. Because some of the rebounding matter is bombarded by the neutrons, some of its nuclei capture them, creating a spectrum of heavier-than-iron material including the radioactive elements up to and likely beyond uranium. Although non-exploding red giants can produce significant quantities of elements heavier than iron using neutrons released in side reactions of earlier nuclear reactions, the abundance of elements heavier than iron (and in particular, of certain isotopes of elements that have multiple stable or long-lived isotopes) produced in such reactions is quite different from that produced in a supernova. Neither abundance alone matches that found in the Solar System, so both supernovae and ejection of elements from red giants are required to explain the observed abundance of heavy elements and their isotopes.

The energy transferred from collapse of the core to rebounding material not only generates heavy elements, but provides for their acceleration well beyond escape velocity, thus causing a Type Ib, Type Ic, or Type II supernova. Note that current understanding of this energy transfer is still not satisfactory; although current computer models of Type Ib, Type Ic, and Type II supernovae account for part of the energy transfer, they are not able to account for enough energy transfer to produce the observed ejection of material.

Some evidence gained from analysis of the mass and orbital parameters of binary neutron stars (which require two such supernovae) hints that the collapse of an oxygen-neon-magnesium core may produce a supernova that differs observably (in ways other than size) from a supernova produced by the collapse of an iron core.

The most massive stars that exist today may be completely destroyed by a supernova with an energy greatly exceeding its gravitational binding energy. This rare event, caused by pair-instability, leaves behind no black hole remnant. In the past history of the universe, some stars were even larger than the largest that exists today, and they would immediately collapse into a black hole at the end of their lives, due to photodisintegration.

Variable Stars

A variable star is a star whose brightness as seen from Earth (its apparent magnitude) fluctuates. This variation may be caused by a change in emitted light or by something partly blocking the light, so variable stars are classified as either:

• Intrinsic variables, whose luminosity actually changes; for example, because the star periodically swells and shrinks.
  • Pulsating variables, stars whose radius alternately expands and contracts as part of their natural evolutionary ageing processes. Pulsating variable stars are classified according to the nature of their variability. Look any of these up in Wikipedia or a good encyclopedia for details: Cepheids and cepheid-like variables (Classical Cepheid variables, Type II Cepheids, RR Lyrae variables, Delta Scuti variables, SX Phoenicis variables, and Rapidly oscillating Ap variables); Long period variables (Mira variables, Semiregular variables, Slow irregular variables); Beta Cephei variables; Slowly pulsating B stars; PV Telescopii variables; RV Tauri variables; Alpha Cygni variables; Gamma Doradus variables; Pulsating white dwarfs; and those with Solar-like oscillations.
  • Eruptive variables, stars who experience eruptions on their surfaces like flares or mass ejections: Eruptive variable stars (Protostars, like Herbig Ae/Be stars, Orion variables, and FU Orionis variables); Giants and supergiants like Luminous blue variables; Yellow hypergiants; R Coronae Borealis variables; Wolf–Rayet variables; Gamma Cassiopeiae variables; Flare stars; and RS Canum Venaticorum variables.
  • Cataclysmic or explosive variables, stars that undergo a cataclysmic change in their properties, like novae, supernovae, dwarf novae and Z Andromedae variables.

• Light curve of eclipsing binary Kepler 16

  Extrinsic variables, whose apparent changes in brightness are due to changes in the amount of their light that can reach Earth; for example, because the star has an orbiting companion that sometimes eclipses it (including Algol, Beta Lyrae and W Ursae Majoris variables, and planetary transits), or because it rotates. Rotating variables (with fast rotation speeds causing them to become ellipsoidal in shape) include non-spherical stars and Ellipsoidal variables; those with extreme “sunspots” (FK Comae Berenices and BY Draconis variable stars); and those whose magnetic fields cause the variation in light seen on Earth (Alpha² Canum Venaticorum and SX Arietis variables and Optically variable pulsars).

These subgroups themselves are further divided into specific types of variable stars that are usually named after their prototype. For example, dwarf novae are designated U Geminorum stars after the first recognized star in the class, U Geminorum.

Many, possibly most, stars have at least some variation in luminosity: the energy output of our Sun, for example, varies by about 0.1% over an 11 year solar cycle.

Variable star nomenclature: In a given constellation, the first variable stars discovered were designated with letters R through Z, e.g. R Andromedae. This system of nomenclature was developed by Friedrich W. Argelander, who gave the first previously unnamed variable in a constellation the letter R, the first letter not used by Bayer. Letters RR through RZ, SS through SZ, up to ZZ are used for the next discoveries, e.g. RR Lyrae. Later discoveries used letters AA through AZ, BB through BZ, and up to QQ through QZ (with J omitted). Once those 334 combinations are exhausted, variables are numbered in order of discovery, starting with the prefix V335 onwards.

R136a1 is a Wolf–Rayet star and the most massive star known, an estimated 265 solar masses. It is also the most luminous star known at 7,400,000 times the luminosity of the Sun. It is a member of R136, a super star cluster near the centre of the 30 Doradus complex (also known as the Tarantula Nebula), in the Large Magellanic Cloud. Zooming in from the Tarantula Nebula to the R136 cluster, with R136a1/2/3 visible as the barely resolved knot at bottom right. The brightest star just to the left of the cluster core is R136c, another extremely massive WN5h star