Supernovae PART 3: Massive Stars

For Supernovae PART 1: Introduction, click here.
For Supernovae PART 2: Low Mass Stars, click here.

MASSIVE STARS: SOME EXPLODE, SOME DON'T

Stars with less than 8 M mass end up as white dwarfs (see PART 2) but those between 8 M and 10 M might explode as supernovae instead. Below is an artists' impression of supernova 1993J, observed 23 years ago. It blew up in Messier 81 Galaxy, 12 million light-years away. The original star was about 10 times more massive than the Sun (10 M) and about 1000 times brighter.

NASA, ESA, and G. Bacon (STScI);Wikipedia

The tipping point here is not completely understood. The star mass range that will explode depends on the model used. While red dwarfs are by far the most common stars in the universe, these more massive 8 - 10 M stars, in turn, are much more common than supermassive stars over 12 M. This means that 8 - 10 M stars probably represent most of our observed supernovae. Those at the lowest mass limit seem to explode as faint II-P supernovae. These are faint relatively low-energy explosions in distant galaxies, so, though common, they are not as easy to observe as larger brighter supernovae.

The current upper limit of stellar mass is at least 150 M and might be up to 250 M, although even more massive stars up to 1000 M might have lived (briefly) when the universe was very young. One might expect a direct relationship between stellar mass and the intensity of the eventual supernova, but the reality is far weirder. In some cases the supernova is unexpectedly faint or absent altogether.

A Word on Supernova Classification

The last two decades, once supernovae could be regularly observed and studied, have revealed an amazing variety of explosions. Attempts to classify all of them is challenging at best. Ultimately researchers would want to classify them based on the star's initial mass, its metallicity and the mechanics of the explosion. Because these parameters take time to figure out, supernovae, when they are first observed, are placed instead into an easier quicker scheme based on the spectra of electromagnetic (EM) radiation they emit. Based on this, five basic classes (Type 1, Type II, Type III, etc.) are often discussed, which are further divided up into subclasses (Type II-P, Type 1a). Instead of an exhaustive survey of them, we'll continue to go upward in star mass, comparing how the stars die and why.

Mass-Exchange Type 1a Supernovae

While supernovae in general vary greatly in terms of their underlying mechanisms, their output energies and their duration, Type 1a supernovae are a unique situation. They are so similar to one another by all accounts that they can be used as standard candles. Their light curves resemble the graph below left where luminosity is plotted against time. All Type 1a supernovae have the same peak luminosity (or absolute brightness). The radioactive decay of nickel (a product of the star's runaway fusion leading up to the explosion, as we will see) creates the peak in brightness. Then cobalt decays, emitting EM radiation.

Xenoforme-commonswiki;Wikipedia

If you know that the absolute brightness of a Type 1a supernova is always the same, you can measure its observed brightness and then calculate its distance from you using the inverse square law, where light from an object decreases at a rate proportional to the inverse square of its distance from an observer. These supernovae occur all over the universe, so they can be used to estimate how far away various distant galaxies are. The discovery of these supernovae was also used to prove that the universe is undergoing accelerating expansion. They are, however, rare -  detected only about once every 100 years on average.

Unlike other supernovae, Type 1a supernovae require more than one star. The explosion mechanism itself is well worked out but how events come about to trigger it is not entirely understood. One thing is certain. Individual stars never explode as Type 1a supernovae. At least a binary pair of stars is required. According to the current model, one of those stars must be a white dwarf. The other star can be any star or a stellar remnant, including another white dwarf, as shown below in this brief European Southern Observatory (ESO) animation.


This is an animation of what is expected to happen to the planetary nebula Henize 2-248 when two white dwarfs orbiting each other at its centre merge in about 700,000 years. The ensuing Type 1a supernova will completely destroy both stars.

The current Type 1a supernova model is based on a carbon/oxygen white dwarf in particular. The composition of white dwarfs can vary from helium (hydrogen fusion inside low-mass red dwarfs) to carbon/oxygen (helium fusion inside Sun-like stars) to oxygen/neon/ magnesium (carbon fusion inside 8 - 10 M stars). Most white dwarfs that exist today are carbon/oxygen white dwarfs.

A slowly rotating carbon/oxygen white dwarf accretes matter from a companion star until it just about overcomes electron degeneracy pressure. This mass limit is called the Chandrasekhar limit (1.44 M). Just as the star is about to reach it, increasing pressure within the star ignites carbon fusion. It will continue to fuse carbon over a period of about 1000 years until the fusion reaction ignites a flame front, which heats the star enough to trigger oxygen fusion. Although the flame front mechanism is up for debate, the star begins to heat up quickly once this point is reached, an event not unlike the helium flash described earlier. In just a few seconds, most of the carbon and oxygen are fused into various heavier elements. The star is degenerate so it can't expand to cool off and regulate its rate of fusion, so it evolves rapidly into a runaway reaction that blows it up in a supernova. The shockwave from the explosion is particularly violent. This is why they can be seen from even extremely distant galaxies. The material blowing up is intensely hot dense electron-degenerate matter undergoing runaway fusion. It is estimated to travel up to 20,000 m/s (compare this to typical supernova shockwave velocities of 40,000 km/h, equivalent to about 11,000 m/s). Because the starting mass is always the same, these supernovae have very similar absolute magnitudes of close to -19.3. That's about five billion times brighter than the Sun. In fact, a Type 1a supernova can outshine an entire galaxy. Although the theory is not worked out, rapidly spinning white dwarfs might be able to exceed the Chandrasekhar limit before going supernova, perhaps into an even brighter and more powerful explosion. Of course, these supernovae would not be typical Type 1a explosions.

Current theory suggests that binaries containing oxygen/neon/ magnesium white dwarfs won't go supernova when the white dwarf accretes mass. Carbon fusion can't be triggered because the star is already full of carbon fusion products. Like carbon/oxygen dwarfs, these stars heat up under increasing pressure, but before higher-level fusion temperatures are reached, the star exceeds the Chandrasekhar limit and collapses further into a neutron star, an even denser state of matter in which the only outward pressure preventing total collapse into a black hole is the strong force.

Although there is not much theoretical knowledge to work from, a helium white dwarf accreting mass in a binary system might not go supernova either. Helium white dwarfs are extremely rare, although they will be common in about a trillion years. No red dwarfs have lived long enough to evolve into one, so observed (all extremely low mass) helium dwarfs are thought to have formed during the evolution of certain binary pairs. Helium fusion could ignite as the helium white dwarf reaches sufficient pressure but it might release that energy in a series of helium flashes rather than an explosion or further collapse. A helium flash is energetic enough to relieve the star of its degeneracy state but not so energetic that it blows the star apart. The star could continue to accrete mass in this fashion until it is massive enough to explode as a mid-mass star core collapse supernova (which will be discussed next).

Type II Supernovae

A Type II supernova, like the faint II-P mentioned, is distinguished by the presence of hydrogen lines in its spectrum, which means that these stars still have a significant outer hydrogen shell when they blow up. Stars between 8 and 50 M typically explode as Type II supernovae. These stars are usually found in the arms of spiral galaxies like our own. They end their lives in rapid core collapse leading to a violent explosion, like the one in the centre of the image below left.

NASA, ESA, P. Challis, and R. Kirshner (Harvard-Smithsonian
Center for Astrophysics);Wikipedia

This Type II supernova called SN 1987A, was one of the brightest supernovae witnessed in modern times. It blazed as bright as 100 million Suns for months after the initial explosion. A pink ring, about 1 light-year across, glows brightly as a shockwave blasts against the ring of material shed by the star approximately 20,000 years before it exploded.

Not all 8 - 50 M supernovae are Type II, however. A number of Type 1 supernovae can also occur. Their spectra have no hydrogen lines. Some massive stars shed their outer shell of hydrogen (Type Ib) and some shed both hydrogen and helium (Type Ic) before their cores collapse and explode. These explosions seem to be limited to mid to high-metal stars that are at least 40 M. Many of these stars, once in the red giant phase, are unstable and they express it in as many ways as we middle-aged people express our mid-life crises. Some shed matter steadily, but others rhythmically pulsate, throwing off layer after layer. The least stable stars flare up in irregular violent bursts, some violent enough to shed their entire hydrogen/helium envelope altogether, leaving them "naked" as they rage into the ensuing supernova. White dwarfs destined to become Type 1a supernovae have no outer layers of hydrogen or helium to start with before they explode. This is what puts them in the Type 1 category.

Faint Type II-P supernovae, those often associated with the smallest mass stars to explode, were mentioned earlier. Not all Type II-P supernovae are faint, however. Some, such as SN 1987A, the explosion of a massive 20 M star, are exceptionally powerful. As a point of interest, SN 1987A was a peculiar Type II-P event. The "P" in Type II-P refers to "plateau." Its EM spectrum more or less maintains its luminosity for a few months after the initial blast, whereas the light curve from a Type II-L (linear) supernova linearly drops off, as shown below right in a simplified graph. The comparison of these two Type II supernovae provide a good example of how light curves tell physicists what's going on during these supernovae. Stars with a wide range of masses can end as Type II-P supernovae. Although all have a plateau, the luminosity of the explosion varies widely, depending on the star's mass.

Xenoforme-commonswiki;Wikipedia

Unlike Type 1a supernovae in which the entire star is a core remnant that heats up and blows itself apart, all Type II, Type 1b and Type 1c supernovae come about when the core starts to collapse. The first sign of a core collapse supernova is a burst of invisible and difficult to detect neutrinos. A few hours later the shockwave itself breaks out of the star, releasing an intense burst of EM radiation, usually an ultraviolet flash. At first the photons can't escape. They are trapped in a thick envelope of ionized hydrogen around the star. Once the hydrogen cools enough to return to its neutral atomic state, the layer turns transparent. At this point the supernova becomes optically visible as it expands. A peak in the visible light curve occurs when the surface area of the star is increasing while its temperature has not had a chance yet to decrease. The time spent in the plateau phase for SNII-P depends on the thickness of the hydrogen shell. A thicker shell means a longer plateau. SNII-L stars have a much smaller hydrogen shell to start so light leaves in a sharper single burst. In both cases, visible light drops off to a radioactive tail, where light is emitted from the conversion of unstable cobalt-56 into stable iron-56.

Before the blast, these stars started to fuse elements heavier than helium in their cores. If the core contracts sufficiently to reach a temperature of about 1.1 GK (gigakelvins or 10⁹ Kelvins), nuclei such as neon, created by carbon fusion, will begin to partially disintegrate into alpha particles (helium nuclei) and gamma radiation. Other neon nuclei capture these alpha particles to create magnesium, while still others absorb gamma photons to create oxygen. Oxygen then fuses to form sulphur, silicon and smaller amounts of various larger elements in the core.

This decomposition process depends on the mass of the star. Only stars of about 8 - 10 M and more undergo this process and blow up as supernovae. While the cores of lower mass stars like the Sun collapse too (into white dwarfs), there is no decomposition and no electron capture, an additional process that is described next. For these low mass stars, the outer layers blow away violently but not explosively, leaving a degenerate inert core behind. If the star is over 8 - 10 M but its core is not massive enough to convert all the neon into oxygen and magnesium, fusion will eventually slow down and it will begin to collapse. The matter in the core by this time is already in an electron-degenerate state. As gravity becomes the dominant force, electron degeneracy pressure can no longer prevent further collapse. Electrons, already squeezed into lowest energy orbitals are now squeezed into the nuclei themselves, a process that is called electron capture. Nuclei capture electrons from their innermost shells. When an electron is absorbed, it changes a proton into a neutron while emitting an electron neutrino (an example of a weak interaction). This process creates smaller amounts of additional elements such as aluminum and sodium in the core. Eventually, the star becomes layered like an onion with elements that fuse at lower temperatures (starting with hydrogen and then helium) occupying the more outermost shells. The diagram below left is a simplified not-to-scale cross-section of a massive star. All the elements are in a plasma state (nuclei are free in a now-diminishing sea of electrons), where the largest nuclei "rain down" through inner levels.

User:Rursus;Wikipedia

Although this diagram, like most, shows an iron core, some supernova-destined stars are not massive enough to fuse iron. The deepest core elements in the least massive supernova stars (those 8 M to 10 M) are more likely to be oxygen, neon and magnesium. In all cases, however, the degenerate core is receiving a continuous injection of energy from gamma radiation as well as an energetic burst of electron neutrinos from the electron capture processes. In this extreme energy environment, the atomic nuclei themselves begin to photodisintegrate. High-energy gamma photons begin to break up the nuclei into alpha particles.

The core is degenerate so it cannot exert fusion pressure and expand to release the heat. Instead, gravity remains dominant. It pushes on the core until it becomes so dense that even neutrinos, for which all atomic matter is normally invisible, are trapped. The additional energy of the trapped neutrinos leads to a massive spike in energy. The core suddenly and violently implodes into itself. What remains of the core is destined to be a neutron star unless the original star was very massive. In that case the core's matter will collapse completely into a stellar black hole. Stars with a wide range of masses end up as neutron stars but all neutron stars are the same diameter, roughly between 10 to 30 km, and the same mass, roughly between 1 and 3 M, depending on the theoretical model used. Think of the Sun's entire mass squeezed into a sphere the size of Earth. That's about a millionth of its original volume, giving you an idea of how much of an atom's volume is empty space and how much it can be squeezed. The rest of the star will be explosively blown away. There are approximately 100 million neutron stars in the Milky Way alone.

In neutron stars, the only force preventing total collapse into a black hole is neutron degeneracy pressure, an expression of the fundamental strong force. As mentioned earlier, the electron degenerate matter of a white dwarf is matter that is prevented from further collapse by the counterforce described by the Pauli exclusion principle. The matter cannot contract or expand in response to changes in temperature, but its electrons can move faster or slower. In the hottest white dwarfs, electrons are so fast they escape their atoms to create a sea of nuclei (mostly alpha particles) within a sea of fast-moving free electrons. There is a limit to how fast the electrons can move and that is the speed of light. As this limit is approached, which is the Chandrasekhar limit attacked from a different angle, electron degeneracy pressure can no longer support the matter. The matter suddenly collapses into neutron degenerate matter.

A Note About Stellar Remnant Mass

There is some discrepancy between models at which core mass will trigger total collapse into a black hole. For white dwarfs, the current often-cited Chandrasekhar mass limit for electron degenerate mass is 1.39 M, although a value of 1.44 M is also often cited. Other models suggest that within very powerful magnetic fields, super-white dwarfs, with masses over 3 M might be stable. As mentioned earlier, rapidly spinning white dwarfs might also have a higher mass limit (and these will likely also exhibit powerful magnetic fields). Generally, however, 1.39 M is the limit for electron degenerate mass. The analogous upper limit of neutron degenerate mass (a neutron star) is called the Tolman-Oppenheimer-Volkoff (TOV) limit. It is between 2 and 3 M. The range in values for both mass limits points out where science continues to be a work in progress. The equations of state used to calculate these masses, especially the TOV mass, are not well understood for matter that is no longer in its ordinary atomic state. No physical lab can supply the kind of energy required to directly observe matter in either degenerate state. For those of you interested in a more detailed exploration of the physics of neutron degenerate matter, try this 2010 paper, intended as an online course.

At around 2 to 3 M, some stellar cores might collapse into a hypothetical intermediate state of exotic quark matter, into a quark star in other words. In this case, neutrons themselves can't hold up intact but matter doesn't completely collapse into a black hole. Instead neutrons break down into a sea of free quarks and gluons, the particles that make up neutrons and protons. The strong force, partially overcome here, normally confines quarks into neutrons and protons. Any further pressure (mass over an upper limit of 3 M) would result in total collapse into a black hole. Black holes in general have an upper mass limit of a gargantuan 10¹⁰ M in theory but, based on observational data, stellar black holes range from five to dozens of solar masses.

If the core mass is less than the TOV limit, neutron degeneracy pressure will prevent further collapse. This pressure acts like a wall during core collapse. The imploding core slams inward, hits neutron degeneracy pressure with the power of the strong force, is immediately stopped in its tracks, and rebounds hard, producing a shockwave that expands outward in all directions. The shockwave explosively expels all the outer stellar material into space, leaving behind a neutron star.

A Mysterious Variety of Neutron Stars

All neutron stars spin very fast and have very powerful magnetic fields, particularly right after they form. The black sphere in the centre of the brief animation below is a rotating neutron star. The lines that curve around it are magnetic field lines and the pink cones emanating from it are EM emission zones.

(Jim Smits;Wikipedia)

Neutron stars emit EM radiation and that is how they are detected. Although neutron stars consist mostly of compact neutrons, the outermost layers, under less intense pressure, are thought to be composed of electrons and protons. The radiation, most often observed as radio waves, is the result of electrons accelerating along powerful magnetic field lines between the magnetic poles of the star and emitting curvature EM radiation. Photons interacting with the magnetic field can create electron-positron pairs that emit additional (gamma) radiation.

The angular momentum of the original spinning massive star is conserved. All stars rotate, and they vary greatly in diameter. If the original star was very large and rotating fast, the much tinier neutron star will rotate very fast. The fastest observed rate is 716 rotations per second. Neutron stars also have very strong magnetic fields and this is mysterious because neutrons are electrically neutral particles. There are several possible explanations for it. First off, as mentioned, most researchers agree that some electrons and protons remain in the uppermost portion of the neutron star's crust where there is insufficient pressure to maintain the matter in a neutron degenerate state. There may be enough of them, and they may be moving fast enough, to maintain the original magnetic field of the star, perhaps acting as a kind of magnetic dynamo. Some experts suspect that the magnetic flux of the field itself is conserved and compressed into the much smaller neutron star. Others suggest that the remaining protons exist in an exotic superconducting state, which can multiply the magnetic field. Still others think that a fossil magnetic field remains frozen in the collapsing plasma that formed the neutron star. You might find this conversation at researchgate.net especially interesting as experts wrestle with this mystery.

Some neutron star remnants can have an even more exceptionally powerful magnetic field. These mysterious variants are called magnetars. Generating the most powerful magnetic field known, as much as 10¹¹ tesla, it would distort the electron clouds in the atoms of your body and kill you instantly from 1000 km away.

Born from a very large star, a highly magnetized neutron star often rotates exceptionally fast as well. Electrons can be accelerated so violently they emit an intense binary jet of EM radiation that, if directed at Earth, identifies it as a pulsar because the radiation appears to rapidly pulse. Both magnetars and pulsars are believed to be very young neutron stars. Eventually their rotation rates wind down as energy is lost to the magnetic field and as EM radiation as well. Even an old neutron star is mysterious. Researchers don't really know if a core of even denser quark matter exists deep inside.

Robert Schulze;Wikpedia

An apple-sized inner core of matter could be squeezed so hard it overcomes quark degeneracy pressure to become even denser electroweak matter. These hypothetical physical states could influence the behaviours of neutron stars.

Next we explore the sometimes unexpected ways in which massive stars of between 10 and 150 M end their lives.