Tuesday, October 18, 2016

Supernovae PART 4: What Happens to Super-Massive Stars?

For Supernovae PART 1: Introduction, click here.
For Supernovae PART 2: Low Mass Stars, click here.
For Supernovae PART 3: Massive Stars, click here.


As we've seen in PART 3, stars with 10 - 50 solar masses tend to end their lives much more violently than less massive faint supernova stars do, but the process is essentially the same. In the cores of these stars, there is enough gravitational pressure for fusion to smoothly continue as neon, oxygen and eventually silicon fuse. If a core temperature of about 3 GK is reached, silicon and sulphur fuse with alpha particles released by the photodisintegration of these and other elements, eventually creating, in a step-wise nucleus-building fusion process, nickel-56. Of all the elements, iron-58 and nickel-62 have the highest binding energy per nucleus. Fusion into the slightly smaller nucleus of nickel-56, produced when iron-52 fuses with an alpha particle, is the first reaction that consumes energy rather than releases it. This means that fusion stops at nickel-56. It is stellar ash. Nuclei larger than nickel-62 release energy when they split apart (nuclear fission) rather than when they fuse. Although the onion diagram we saw in PART 3 shows an innermost core of iron, nickel-56 is the last fusion product in any stellar core, regardless of stellar mass. Any further fusion consumes energy so fusion abruptly stops and the core collapses. Meteorites and rocky planets contain significant iron-56. This is the radioactive decay product of unstable nickel-56.

In 10 - 50 M stars, like all stars, the outward pressure of nuclear fusion initially keeps matter in its ordinary atomic state (the main-sequence phase). As fuel is consumed and fusion eventually slows down, matter is squeezed into an electron degenerate state. Then, electron degeneracy pressure can no longer hold up the core, which exceeds the Chandrasekhar limit, a core mass of approximately 1.4 M. By this point, the star is on its way to an iron-core collapse. Electron degeneracy is overcome and the inner core implodes explosively, over a course of just a few seconds. The outer core follows inward, reaching a velocity of almost  the speed of light. The core temperature spikes and electron capture proceeds rapidly, transforming the inner core into a neutron star. Neutron degeneracy pressure acts like a wall, as mentioned, creating a shock wave that redirects the implosion outward. In stars in this range the shockwave is more intense, resulting in a typical Type II-P supernova. Most of this energy comes from a 10-second blast of thermal neutrinos. These particles are much more abundant than the electron capture neutrinos formed earlier. Thermal neutrinos are formed at about 100 billion K as part of the pair production of neutrino/antineutrino pairs (in all flavours). With a whopping combined energy of about 1046 joules, about 10% of the star's mass is carried off in a thermal neutrino blast. This is the explosive energy of the supernova. To get an idea of just how much energy 1044 joules is, check out this energy scale page and keep scrolling down. Of course, neutrinos are invisible. Through some process that isn't yet understood, most of the blast energy (1044 joules) must be reabsorbed into the core somehow to produce the intensely bright EM explosion (a lot of which is in the visible range). During the brief microseconds before pressure and temperature dissipate, a cornucopia of elements heavier than iron, including heavy neutron-rich radioactive elements, are fused and become part of the expanding cloud.

Origins of the elements are shown in the helpful periodic table below.

For stars between 40 M and 90 M, some outer material might fall back onto the newly formed neutron star. If the star's core grows massive enough, estimated at between 2 and 3 M, it will overcome neutron degeneracy pressure (the TOV limit) and collapse into a black hole. The inward fallback of matter can reduce the outward kinetic energy of the explosion, so much so that in some cases there is no supernova at all. In some cases, the in-falling matter (particularly if it is rapidly spinning) generates two opposing jets of matter traveling outward at close to light speed, a brief event called a gamma ray burst (GRB) that can last from a few milliseconds to hours, or an exceptionally bright supernova (called a hypernova) can result, or both.

A massive star's mass, along with its metallicity, determines its eventual fate whether it will become a neutron star, black hole or leave behind nothing at all. Some stars will even collapse into a black hole with no accompanying supernova. The screen shot below from Wikipedia's Supernova entry lists possible stellar fates.

This chart showcases the myriad ways in which massive stars can die. How a star dies depends primarily on two factors: how massive it is and of what it's made. Low-metal stars contain very few atoms more massive than helium as they start out on the main sequence of their lives; high-metal stars have a larger but still tiny percentage of heavier atomic nuclei. An enormous range of stars from 8 M up to 250 M can explode, and that explosion can be faint, typical or tremendously bright.

A Note on Gamma Ray Bursts

A gamma ray burst (GRB) is the most energetic event in the universe outside of the Big Bang itself. The most powerful one ever detected was the 2008 GRB (called GRB 080916C). It was 2.5 million times brighter than the brightest supernova ever detected. To offer perspective, a typical Type II-P supernova usually outshines an entire galaxy. For a few seconds in 2008, its light, from 7.8 billion light-years away, was visible even to the naked eye.

Gamma ray bursts appear to accompany the collapses of very massive rapidly spinning stars. They were first discovered in the 1960's when the U.S. detected what they suspected were flashes of radiation from secret Soviet nuclear tests. Researchers began to realize that these randomly located, powerful transient bursts did not accompany any obvious objects or stars in the Milky Way. In fact, none have ever been detected in our galaxy, though they may have occurred in its past. Through a process not well understood, much of the energy (and a significant amount of the mass) of a collapsing massive star converts into gamma radiation released in two tightly focused bilateral jets that can travel across the universe. The 2008 GRB described had an estimated energy output of 9 x 1047 J, about five times more energy than the equivalent of the entire Sun's mass.

Sometimes described as a shot heard clear across the universe, these extremely energetic photons start out as gamma rays and stretch as the universe expands while they travel. Billions of years later, they are detected as much longer, less energetic, but still intense X-ray, radio or even infrared bursts. Although GRB's are fairly frequent events in the universe at large (about one per day is detected), Earth's detectors don't pick up the vast majority of them because we must be lined up with the narrow jet of the GRB in order to detect it.


Pair Instability Supernova

Very massive stars (over 140 M) may be completely destroyed by their powerful supernovae. A process called pair instability leaves nothing, not even a black hole, behind. In the core of such a star, collisions between atomic nuclei and gamma photons are so violent that the new gamma photons (shown as white squiggly lines below) created from those collisions produce new matter/antimatter pairs - electrons and positrons (black and white spheres below) - via a process called pair production. The energy of the gamma photon must be extreme, equivalent to the rest mass of the particle pair (0.511 MeV x 2) in order for this to happen. Each particle pair created drains significant energy from the core as energy is converted into matter. As energy decreases, outward pressure drops and the core begins to collapse.

NASA/CXC/M. Weiss;Wikipedia
The temperature in the massive core is so high by the time it starts this final contraction that runaway fusion reactions generate enough energy to blow entire star into space, sometimes accompanied with a particularly bright GRB. Only a nebula of ionized gases and heavy elements mark the original location of the star.

Pair production reenacts matter creation that occurred just after the Big Bang. At a threshold temperature of about 10 GK, gamma photons convert back and forth into electron/positron pairs in an equilibrium state between energy and matter. Temperature is the average kinetic energy of particles so some gamma rays will be more energetic than average and initiate pair production well before this temperature is reached in the core. There is tremendous core energy in all extremely massive stars, but only stars within a specific stellar mass and metallicity range can undergo pair instability.

Stars below 100 M aren't massive enough to trigger core pair production. They follow one of explosion pathways typical of stars between 40 M and 90 M. Stars between 100 and 140 M are large enough to trigger some pair production but it isn't energetic enough. There aren't enough pairs produced and photons taken out of the core to reduce the outward pressure enough to trigger runaway fusion. Instead the fusion rate is increased just enough to return the star to equilibrium. These stars go through several increased fusion/pair production cycles in a series of pulses, each time losing stellar gas, until the star mass falls below 100 M. Stars between about 140 and 250 M are true pair instability stars. In these stars, more thermonuclear energy is released than the entire star's gravitational binding energy. The entire star is ripped apart in what can be an exceptionally powerful supernova (a hypernova). Not all of its matter is converted into radiation, however. Up to dozens of solar masses of unstable nickel-56 can be blown away from the core, which decays into cobalt-56 and then into stable iron-56, contributing to an iron-rich stellar nebula. Solar material can also be sprayed across the universe in a powerful GRB as well. The radiation contributed by various heavy nuclei decay reactions could make such a supernova exceptionally luminous and long-lived as well.

Stars Over 300 M

The most massive known star is R136a1. Shown in the center of the 2010 near-infrared image below, Hubble Space Telescope resolved it from the R136 concentration of stars, located 156,000 light years away, in 1992.

ESO/P.Crowther/C.J. Evans;Wikipedia
There are many massive stars in this cluster. R136a1 is the most massive, hottest and most luminous of any known star. It would occupy the extreme top left corner of the Hertzsprung-Russell diagram. Formed from an initial mass of 325 M, it has a current mass of 315 M. Less than a million years old, intense EM radiation from this very hot star creates powerful solar winds that strip away its mass at a rate a billion times faster than the Sun. This star is close to the Eddington limit. Calculated using hydrogen plasma, stellar winds from a star any more luminous would quickly strip away its mass before it could evolve further. GRB's and supernovae, by the way, greatly surpass the Eddington limit and that is why they are so brief and mass loss is so intense. R136a1 is still fusing hydrogen in its core, using the CNO cycle because its core is very hot. Significant levels of helium and nitrogen in its spectrum reveal that it is also strongly convective.

Expected to stay on the main sequence for only 1.7 million years, stars such as this live hot and die young. By the time it uses up its hydrogen and evolves into a luminous blue variable star, it will have only about 80 M left. As carbon and oxygen accumulate in its helium core, core temperature will continue to rise and mass loss will increase further. After several hundreds of thousands of years the helium will be used up. Heavier element fusion will last only a few thousand years. Though estimates vary greatly its massive core will collapse and likely trigger a supernova that may or may not leave behind a black hole. Its supernova spectrum will likely classify it as a Type 1c explosion because by the time it explodes, its hydrogen and helium will have long been blown away by stellar winds.

It is expected to have a metallicity similar to the Large Magellanic Cloud in which it exists, about 1/4 that of the Sun. The Large Magellanic Cloud is a gas-rich metal-poor neighbor galaxy full of star-forming nebulae and young populations of stars. Below, enormous R136a1 (bright medium blue) is compared a 0.1 M red dwarf, the Sun (shown incorrectly yellow) and an 8 M blue star (pale blue), all in the main-sequence phase of their lives.

ESO/M. Kornmesser;Wikipedia
Massive Metal-free Stars

Despite the confirmation of R136a, stars over 300 M remain a bit of a mystery. Estimates of how these stars die vary greatly depending on the model used. While it is possible that they could be totally destroyed in pair-instability supernovae like slightly less massive stars, it is also possible that they could end as photodisintegration, rather than pair production, triggers core collapse. This process can be modeled with a massive metal-free star, consisting of hydrogen, helium and a tiny percentage of lithium and beryllium - precisely the kind of first stars to form from the pristine gases and dust left over from the Big Bang. These stars fused the first heavier elements to exist in the universe.

Unlike R136a1, no zero-metallicity stars, stars formed strictly from primordial material left over form the Big Bang, have been directly detected, but theory points to their once existence. Because we look back in time as we peer across the universe, these stars, if they exist, will be located extremely far away at the extreme edge of the visible universe. They may have shone just as the first light from the Big Bang itself was able to stream through space, and space itself would have been just 1/30th of its present size. This would result in extremely red-shifted extremely dim light from once exceptionally bright bluish white behemoths.

This might be what those stars look like today. Not quite what one might expect at first thought.

This infrared image taken by the Spitzer Space Telescope in 2005 has all the stars, galaxies and artifacts greyed out and the background enhanced to reveal a glow not attributed to present stars or galaxies and not attributed to cosmic background radiation. Researchers are uncertain but it might be the extremely red-shifted light from the first stars to shine in the universe.

Below, an artist's impression shows what those first stars may have looked like at the time, just 400 million years after the Big Bang. Their light would have been the first light to shine from any object ever.

NASA/WMAP Science team
It is possible that such distinctly metal-free stars would not be visible at all to us today. If these stars were convective, they could dredge up their fusion products from their cores to their surfaces, making them indistinguishable from regular low-metal stars. These stars could have lived their short lives before EM radiation could escape (during the so-called the "dark ages" prior to the recombination period). Interstellar space then was hot dense plasma consisting  of photons, electrons and protons. A process called Thomson scattering made it opaque to all EM radiation. Until the universe expanded and cooled enough for protons and electrons to recombine into neutral atoms, no light, including that from any embedded stars, could shine outward to be detected by us.

Unlike the catastrophic runaway fusion reactions triggered during pair-instability, photodisintegration is an endothermic process for nuclei up to iron (the most tightly bound nucleus). It absorbs energy. And unlike the incomplete photodisintegration process in less massive stars, which knocks off one or two nuclear protons or neutrons, the disintegration in this case should be complete, down to alpha particles and protons. Based on computer modeling, this process should start at about 6 x 109 K.

Most, but not all, modeling suggests that these no-metal stars could have been very massive, perhaps up to 1000 M. Star-forming molecular clouds then would have been much warmer, up to 800 K (525°C) as opposed to cold molecular clouds just a few degrees above absolute zero today. This means that in such an energetic environment a much larger minimum mass would be required to form a star and a much higher Eddington limit could be achieved. These stars would have appeared much like R136a1 - very massive, very hot, very bright and very short lived, burning through their fuel in just a few million years.

Because these stars have no or almost no carbon in them to start with, there would be no carbon to trigger the CNO (carbon/nitrogen/oxygen) fusion cycle. The only fusion reaction available during the main-sequence phase would be the p-p chain reaction of hydrogen fusion. This reaction is much less temperature sensitive than CNO fusion, and this means that it doesn't serve as the built-in thermostat that CNO fusion does. The star's core could therefore get much hotter than it would in stars with metallicity. Eventually it would get hot enough to trigger the triple alpha fusion process and from there fusion would proceed, as carbon is fused from beryllium and helium. Once sufficient carbon is present, the CNO cycle would be triggered as well.

Though fusion would not be runaway, it would be at such a high rate under such enormous gravitational pressure that it would eventually trigger photodisintegration. Gamma photons would be powerful enough to rip apart even the smallest most tightly bound nuclei. In a matter of perhaps minutes, millions of years of core fusion would come undone as gamma photons, more energetic than those in pair instability cores, are absorbed by the core nuclei, causing them to split into alpha particles, free neutrons and free protons. The sudden absorption of free energy (the gamma photons) would trigger catastrophic complete collapse before the core had any opportunity to re-ignite fusion. The core, as a result, would continue to collapse into a massive black hole, with no counter-process to stop it.

If massive low to no-metal stars ended their lives because of photodisintegration, there should be a number of massive black holes left from their destruction. If little or no explosion occurred after the stars collapsed, the black holes should have masses close to the original stars. Unexpectedly bright galaxies (suspected to contain clumps of very massive hot stars) from the period when light first began to shine in the universe are now being detected using a variety of telescopes. Some individual star candidates between 250 and 1000 M have also been indirectly detected. This leaves the mystery of how the young universe seeded itself with nuclei larger than beryllium. A near complete collapse into a massive black hole takes any metals formed right along with it. It is possible that the first stars spewed out their matter in powerful GRB's as they collapsed into black holes (and these most ancient GRB's should be detectable). It is also possible that the first stars were slightly less massive than 250 M, and they might have blown up completely as extremely powerful supernovae instead, spewing out lots of fusion products and leaving behind no trace other than a large radioactive cloud destined to become a star nursery for stars with metallicity.

Despite gas-rich low-metal regions that are still actively making huge R136a-like stars, ancient massive, luminous, hot, metal-free stars could represent the ancient beginnings of a global evolutionary trend toward decreasing stellar mass, as star nurseries in general cool and grow richer in massive elements. The most abundant type of star today is the red dwarf. These stars are so long-lived that if they formed early on in the universe they should still be burning. However, no metal-free red dwarf stars have been observed (although they are dim and small so they would be very difficult to detect at such a great distance). All observed red dwarfs have metal content and almost all of them contain a significant percentage of larger elements, indicating that they were formed in molecular clouds thoroughly seeded with the residues of many previous supernovae. This, along with the fact they are found only in the spiral arms of young galaxies like our own, indicates that red dwarfs tend to be more recently formed than distant and long-dead metal-poor massive stars.


Stars dotting the night sky might seem eternal, but if we could live for millions of years we would bear witness to countless dramatic life and death events. After lives that last anywhere from millions to trillions of years, stars die. Some die peacefully and gradually while others blow up with unimaginable force. Understanding what happens to matter under such atom-ripping conditions is a huge challenge because these forces cannot be replicated on Earth. Gravity crushes atoms in the cores of massive stars into mysteriously dense physical states where the rules of atomic behaviour no longer apply. During some supernovae, matter is crushed completely into a black hole where even the rules of space and time no longer apply. In both cases, new larger atomic nuclei are also created. Scientists are just now able to fuse the largest of these, such as ununoctium (which decays within a millisecond), in the most powerful particle accelerators. It is a job requiring extreme energy, one that is done naturally in microseconds during a supernova.

The universe burst into being almost 14 billion years ago, laced with only the lightest elements, mostly hydrogen along with helium and a tiny amount of lithium that were fused before expansion and cooling halted the process. Gravity shaped that primordial material into stellar fusion reactors that supply the conditions to fuse a far greater variety of elements, up to iron. The synthesis of even heavier elements requires enough energy to overcome a highly endothermic process that absorbs energy rather than releases it. Only in the very brief, chaotic and extremely energetic environment of a supernova can elements such zinc, silver, gold, mercury and lead, as well as heavy unstable elements such as radium, uranium and plutonium come into existence. Furthermore, the explosive process blasts out not only those elements fused in the supernova but those fused earlier inside the star as well. That debris cools into molecular clouds that later become new stars and their planets. Rocky planets like Earth can only form from molecular clouds laced with elements forged in the bellies of massive stars and forged in their violent deaths.

Monday, October 17, 2016

Supernovae PART 3: Massive Stars

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


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.

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);
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.

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.

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.

Thursday, October 13, 2016

Supernovae PART 2: Low Mass Stars

For Supernovae PART 1: Introduction, click here.

Brown Dwarf Stars

The smallest mass that can form into a star is a brown dwarf.

A brown dwarf forms from a protostar mass less than 0.08 M (M = one solar mass or 2 x 1030 kg). This mass is on the edge between being a star and being a planet.

What distinguishes a brown dwarf as a star is its formation. Unlike a planet formed from the protoplanetary disk material around a star, a brown dwarf is formed in the manner described in PART 1. In this case, however, the star's core is too small to have sufficient pressure and heat to ignite the fusion of hydrogen into helium. Fusion does ignite, though. If the mass is at least 0.0125 M (about 13 times the mass of Jupiter), it will trigger the fusion of deuterium into helium, at least temporarily. Deuterium is a stable less abundant isotope of hydrogen that contains a proton as well as a neutron in its nucleus. This fusion reaction requires a lower ignition temperature - roughly 1 x 106 K compared to 4 x 106 K required for hydrogen fusion. This means that 0.0125 M is the smallest possible protostar mass that can evolve into a main-sequence star. Even a large-end brown dwarf will only shine very briefly and dimly. It will eventually cool and die gradually and peacefully into an inert mass. At first only theoretical objects, since the 1990's hundreds of brown dwarfs have been identified by infrared surveys. Almost impossible to see in visible light, the cores of brown dwarfs are compressed enough to emit heat, and that heat gives them away even long after deuterium fusion has fizzled out.

Red Dwarf Stars

A slightly higher mass protostar evolves into a red dwarf star. With a mass between about 0.08 M and 0.50 M, these stars are massive enough to fuse hydrogen into helium through the proton-proton chain reaction. An artist's conception of a typical red dwarf star is shown below right. Though called "red" dwarfs, the surface temperature of these stars means that they would look orange at close range.

Courtesy NASA
Red dwarfs with a mass less than 0.35 M are fully convective, and this makes them uniquely long-lived stars. The helium produced during fusion mixes throughout the star rather than building up in the core as it does with most larger stars. This means that a red dwarf makes thorough use of its hydrogen fuel. It evolves very slowly and maintains a constant luminosity for up to trillions of years before it leaves the main-sequence stage of its life. In main-sequence phase these stars are very dim, emitting between 1/10,000th up to 1/10th of the Sun's luminosity. Although red dwarfs are by far the most common stars in the universe, making up about 75% of all stars, they are difficult to observe. Only very young red dwarfs currently exist. After all, the universe is only 13.8 billion years old. Eventually a red dwarf will fuse all of its hydrogen into helium, end its main-sequence phase, and evolve into a helium white dwarf.

White Dwarf Stellar Remnants

A white dwarf is a stellar remnant. These stars cannot form directly in a nebular cloud. They are denser than any main-sequence star because they are composed of atomic nuclei packed together in a sea of free electrons rather than ordinary atomic matter. Intense force is needed to squeeze atomic matter into such a state. About 97% of the stars in the Milky Way, all stars too small to end up as neutron stars or black holes - from red dwarfs to stars ten times more massive than our Sun - will end up as white dwarfs. When a red dwarf burns up most of its hydrogen fuel, fusion begins to taper off. As the outward pressure from fusion decreases, the star (almost all helium by this time) is overcome by gravitational pressure. Electrons are forced down into orbitals closer and closer to the nuclei, similar to the kind of compression that happens when matter cools off but this case is unique. Even though the matter is compressed, thermal (outward) pressure decreases rather than increases as it normally would because atoms can't move around freely. This means that the core continues to contract. There is a limit to how close electrons can get, however. More than one electron cannot occupy the same quantum state in matter, according to the Pauli exclusion principle. This generates a counter pressure against further collapse. It is called electron degeneracy pressure because, as electrons are forced into lowest possible orbitals, they themselves move faster and faster, generating their own pressure. The helium nuclei at this point can no longer hold onto their electrons so a sea of free fast electrons forms, embedded with helium nuclei. Because it is not in an ordinary atomic state, electron degenerate matter cannot cool off like ordinary matter does. The star will continue to shine white-hot for billions of years and it will only very gradually cool to an inert black dwarf. Heat slowly diffuses outward from its degenerate inner core to a thin outer layer of non-degenerate atomic matter. This thin layer of matter is able to radiate the heat into space, emitted as light, first blue-white, then yellow, orange, and finally red. The brief animation below shows what this process would look like over time.

The heat loss is very inefficient and that is why white dwarfs cool very slowly. Although white dwarf stars on their own die very slowly and peacefully like this, white dwarf relationships with other stars or stellar remnants end badly in explosions called Type 1a Supernovae. They will be explored in PART 3.


Protostars with masses between 0.50 M and 10 M are mid-size stars. These stars are destined to undergo a red giant phase at the end of their main-sequence phase, before they end as white dwarfs. Unlike smaller red dwarf stars, these stars are not significantly convective, The plasma doesn't mix much so a shell of hydrogen remains unburned around the core as it undergoes fusion. Smaller mid-size stars like the Sun utilize the proton-proton chain reaction to fuse hydrogen into helium. With increasing stellar mass, another fusion chain reaction called the CNO cycle becomes the more efficient reaction, fusing hydrogen into helium at higher core temperatures. Whereas proton-proton chain fusion ignites at about 4 x 106 K, the CNO fusion reaction becomes self-maintaining at about 15 x 106 K. There are various reaction paths available, where carbon, oxygen and nitrogen (products of helium fusion) function as catalysts. Stars over 1.8 M almost exclusively utilize to the CNO cycle.

Mid-size stars go through a more complex evolution than smaller red dwarfs. During any star's main-sequence phase, the core is a spherical fusion reaction that maintains hydrostatic equilibrium between outward radiation pressure and inward gravitational pressure. The smaller the star's mass, the longer it will stay in equilibrium as a main-sequence star. Fusion reactions take place at a slower pace in general and increased convection stirs new reactants into the core. The Hertzsprung-Russell diagram (below) plots stars based on luminosity versus surface temperature. This diagram, created in 1910, led to major theoretical developments in stellar physics, well before stellar fusion was understood to be the reaction mechanism in stars (discovered in the 1930's by Hans Bethe). The diagram has evolved since then as theory evolved. It is still one of the most useful tools astronomers have to study stellar evolution. It highlights a number of important trends, the most obvious one being that luminosity tends to increase with surface temperature. You can also see that a star's colour is an indication of its surface temperature. A star's mass determines its position on the main sequence. More massive stars are hotter and brighter than less massive stars. Though shown incorrectly as pale yellow in this case, the Sun is a white (white-hot) star of luminosity 1 (solar unit). It is located in the middle of the down-sloping central spine of the main sequence.

Image created by the European Southern Observatory; original source here.
Our Sun, at 1 M (solar mass), is an example of a mid-size star. The photo below left shows you what the Sun looks like in visible light through a filter.

Geoff Elston; Society for Popular Astronomy:Wikipedia
Its position is currently on the main sequence but it will move across the diagram as it evolves. The Sun will exist as a main-sequence star, fusing hydrogen into helium for a total of 10 billion years. In another 4.5 billion years, its core's hydrogen fuel will be almost entirely fused. At that point, like a victim shot in the saloon of an over-acted Western, the Sun will begin a dramatic and complex series of death throes. The core will start to collapse under its own weight as dwindling outward pressure from hydrogen fusion no longer sustains its volume. The core will heat up under increasing pressure and eventually this process will cause the Sun to expand into a red giant (see the bubble attached to the right of the main sequence in the Hertzsprung-Russell diagram). This change happens because the core has gotten so hot that hydrogen in the shell around the core starts fuse causing the shell to expand. The Sun's surface temperature will cool from white to reddish-orange as it balloons outward to a radius about 200 times its current size (the H-R diagram is not to scale). Despite its cooler surface temperature, it will be almost 3000 times more luminous than it is now because it will be so huge. Its luminosity will increase as it ascends the short red giant branch as the rate of hydrogen fusion in the outer core shell layer increases. The diagram below puts the change in size (200X) into perspective.

Oona Räisänen (User:Mysid), User:Mrsanitazier;Wikipedia 
Having left its main-sequence phase, the Sun will arrive at a new but shorter-lived equilibrium, perhaps lasting hundreds of million of years. It will now occupy the tip of the red giant branch (RGB), the bubble on the Hertzsprung-Russell (H-R) diagram above. As the Sun continues to fuse hydrogen in the thick shell around its core via the CNO cycle, the helium in the core builds up mass. During the last million years or so of the RGB phase, the Sun will lose about 20% of its mass as rapid hydrogen fusion ejects gas from the outermost layers. As hydrogen is used up and fusion slows down, gravity dominates over fusion pressure. Core pressure meanwhile builds as new helium is deposited around it. Gravity squeezes the helium atoms into an electron degenerate state. In stars up to about 2 M, the RGB phase will end abruptly when the inert degenerate helium core reaches the ignition temperature (about 108 K) for fusion into carbon and oxygen using the triple alpha process. Much of the core helium will fuse simultaneously in a violent but invisible process called a helium flash. It will all take only a few minutes because the core is mostly degenerate. Degenerate matter has no opportunity to expand and dissipate the fusion heat. The reaction rate increases into a runaway state, in which a positive feedback loop further increases the reaction rate. All this takes place in an electron sea that is a perfect conductor of heat. The electrons transfer the energy of the fusion reaction throughout the core almost at once, so runaway fusion takes place across it simultaneously. If light could stream away from the Sun at this point, it would exhibit a flash so bright it would be about 1011 times brighter than it is now, about as bright as the entire Milky Way, but it won't be able to. Energetic photons produced by the flash will be absorbed by the Sun's now extremely thick and dense outer envelope of plasma instead.

The force of the run-away helium flash will blast the core nuclei and electrons apart, allowing them to reorganize as ordinary atoms. The core will quickly expand to dissipate the reaction heat. Triple alpha fusion will continue at a now steady rate as the Sun enters an even briefer third stable period, lasting about 100 million years, on the horizontal branch. This branch is not shown in the H-R diagram above. If you scroll downward you will see an evolution graph specific for the Sun. The horizontal branch is a dotted line band. By this time, about 40% of the Sun's total mass is helium and 6% is carbon.

When the helium starts to dwindle, the core will contract further under gravitational pressure and the Sun will expand once again as it resorts to fusing the remaining hydrogen and helium in its outer layers. It will now reach a new stability as an asymptotic giant star, reddish orange, and about 2000 times more luminous than it is now. In another 30 million years, it will begin to eject its outer layers of mass through a series of thermal pulses, each one lasting about 100,000 years, where the core temperature wobbles slightly, causing significant changes in the fusion rate. The ejected outer envelopes of gases will form a planetary nebula around a core remnant that has recompressed into the degenerate matter of a carbon/oxygen white dwarf, containing about half the Sun's current mass squeezed down to the size of Earth. Starting out bluish white hot and about 100 times more luminous than the Sun is now, it will very, very gradually fade into a black dwarf.

The life of the Sun is illustrated in more general terms below, where changes in size are shown to scale. The Sun started out as a relatively small star formed from a molecular cloud. It will spend most of its life as a stable white star until its core begins to run out of hydrogen fuel. It will expand to 200 times its former size into a red giant star. Eventually its outer layers of material will blow away leaving behind an Earth-sized white dwarf (tiny white dot) in the center of a planetary nebula.

ESO/S. Steinhofel;Wikipedia
The graph below tracks the evolution of a 1M star like the Sun.

Though complex and dramatic, the death of a star like our Sun is a mild event by astronomical standards. When our Sun leaves the main sequence, planetary orbits will probably be disturbed by powerful solar wind gusts and the innermost rocky planets will be engulfed as the Sun expands to a red giant. Still, there is some evidence that Earth, though scorched, sterilized and possibly engulfed by hot plasma at some point, might survive physically intact to eventually orbit the Sun's white dwarf remnant.

For all mid-sized stars the evolutionary scenario is similar with some exceptions. Stars slightly smaller than our Sun also evolve into red giants but their cores don't have enough mass to ignite helium fusion. They don't reach the tip of the red giant branch shown above, and there is no helium flash. Instead, these stars move directly into a state that more resembles the post- AGB (asymptotic giant branch) top horizontal line in the graph above, except that they are less luminous. Like the Sun, they eventually blow off their outer layers to form a planetary nebula, leaving behind, in this case, a helium white dwarf. In stars more massive than our Sun, between 2.25 M and 10 M, the core is large enough and dense enough to reach the fusion ignition temperature of helium before it has a chance to contract into a degenerate state. This eliminates the helium flash from their life cycle for a different reason. Like the Sun, these more massive stars expand into red giants as the outer layers of the star cool and become opaque to radiation. However, those layers of gases are even thicker in these more massive stars. They effectively hold in the heat and cause the hydrogen shell around the core to further heat up and expand to form much larger stars, often called super-AGB stars, which tend to evolve further along the horizontal branch than Sun-like stars do. They can become very unstable pulsating stars that can shine yellow, white or even blue-hot. The particular evolutionary path of the star depends on its heavy element content (called metallicity in astronomy) and helium content to start with. The contribution of these elements, in turn, depends on the composition of the molecular cloud in which the star formed. The lines are dotted in the graph above because these horizontal evolutionary paths are still being modeled. The cores of these more massive stars eventually become degenerate and hot enough to fuse carbon into larger nuclei. Some of these stars might ignite a carbon flash analogous to the helium flash of smaller stars. If the original star is between 8 M and 10 M, it will most likely form an initially blue-hot white dwarf surrounded by a planetary nebula, much like less massive mid-range stars, with the exception that the white dwarf will be made of oxygen, neon and magnesium, all of which are products of carbon fusion.

As we've seen here, low mass (small and medium sized) stars die in a step-by-step process that, though violent, is not explosive. At around 8 M in some cases, stars have enough mass to go out spectacularly and these supernovae can vary considerably. Some massive stars buck the trend and don't explode at at all. It all depends on the high-energy physics of matter deep within their cores, next.

Tuesday, October 11, 2016

Supernovae PART 1: Introduction

The Milky Way is full of stars, gas and dust. The bright white dot in the centre of the photo below is Jupiter. The red laser from one of the three observatories in the photo points directly at the heart of our galaxy.


A false-colour image of the Sun taken in the
extreme ultraviolet range of the electromagnetic
spectrum taken by NASA/SDO (AIA)
The photo above of the Milky Way shows us only a small percentage of all stars in the night sky above us. Most of them are invisible to our naked eye from Earth, even those within our own Milky Way.

Stars, like our Sun (right), are gigantic spherical nuclear fusion reactions.

For most of a star's life, the outward blast of nuclear fusion is balanced by the force of gravity squeezing down on the star and it burns steadily. That balance eventually ends when the star runs out of fuel, and when it does it can (but not always!) produce the most violent explosion observed in the universe, a supernova. The Crab Nebula, below, is a six light-year wide remnant of a star that once existed about 6500 light-years away. It exploded in 1054. Not visible is a stellar remnant called a neutron star in the centre. Its powerful magnetic field whips up electrons in the cloud around it almost to the speed of light. They emit the eerie bluish glow.

The image of the Crab Nebula is a mo taken by the Hubble SpaceTelescope. The various glowing colours are different excited ionized elements blown out in the blast.

On Earth, we observe a distant supernova as the sudden appearance of a bright "new" star that fades over a period of weeks or months. A massive star can explode and then disappear into a black hole. Or it can leave absolutely no trace whatsoever of its existence. Often, a remnant remains after the original star dies, composed of the strangest densest matter known. By exploring how stars live and die, we are exploring how matter works at its limits and beyond its limits.

Like us, stars have life cycles. They are born, they change throughout their lives, and eventually they die. While some stars simply fade away, others go out spectacularly as supernovae. The science behind star death is evolving quickly thanks to new robotic high-resolution telescopes that can quickly cover large regions of the night sky and which can recognize even very minute changes in stellar luminosity. Supernovae happen suddenly and fade quickly. Now astronomers can catch one as it unfolds. By continuously scanning across hundreds of galaxies every 30 minutes, the Kepler Space Telescope designed primarily to detect extrasolar planets, for example, can also catch the first minutes of a supernova. In 2011 it caught the initial shockwaves of two brilliant distant supernovae, as two massive red supergiant stars exploded in separate incidents. Watch this brief animation of a supernova shockwave flash. In reality the flash lasts for about an hour. The animation, based on Kepler's observations, was created by NASA's Ames Research Center.

Supernovae are not only fascinating in themselves. They are essential to the creation of all the planets, moons, asteroids and life in the universe. While stars fuse together elements from neon to nickel in their cores, elements with nuclei larger than nickel can only be created in the intense furnace of a supernova explosion. By understanding the physics of supernovae explosions, physicists can understand how they seeded the universe with these heavy elements over time. Using mathematical computer modeling and observational data, they are discovering a surprising myriad of ways in which stars violently end their lives, some of which test the limits of current theory. As the most energetic events in the universe, they offer clues about how the universe is evolving over time.


All stars start out the same basic way. A star's life begins when a cloud of dust and gas in space collapses under its own gravitational attraction. These clouds are called nebulae. A nebula consists of various molecules, neutral atoms of hydrogen and helium, and ionized gases. It can be vast, up to millions of light-years in diameter.

This iconic composite photograph taken by the Hubble Space Telescope is of the "Pillars of Creation," part of the Eagle Nebula. These pillars, composed mostly of molecular hydrogen gas and dust, are star nurseries.

This mosaic photograph of stunning nebulae shows off the Spitzer Space Telescope's capability. These infrared images were used for its 12th Anniversary calendar. As beautiful as they are, nebulae don't last.

Nebulae are transient structures in the universe. The Eagle Nebula, which is about 7000 light-years from Earth (shown below), may not even exist anymore. You can see the Pillars of Creation in the bright white center of the much larger Eagle Nebula below.

This is a three-colour mosaic image of Eagle Nebula taken by the Wide-Field Imager at La Silla Observatory in Chile.

The Spitzer Space Telescope recently imaged a rapidly expanding cloud of hot dust in this region of space. It might be the signature of an intense shockwave produced by a nearby supernova. Knowing the velocity of the shock wave, researchers estimated it will reach the nebula in about 1000 years. We will be able to witness that carnage from Earth about 1000 years from now, not because we will be seeing it in real time but because we are seeing the nebula from 7000 light-years away. When we look up at space, we are looking back in time. We see the Pillars as they were 7000 years ago. If the hot cloud is indeed a shockwave, the nebula was actually destroyed 6000 years ago.

A Refresher Note On Space, Time and Light-speed

I find it is always useful to mentally refresh myself about how space works because it can seem counter-intuitive. ALL electromagnetic (EM) radiation travels at light speed. This includes visible light, gamma rays, X-rays, radio waves, etc. Distant supernovae we observe from Earth exploded hundreds of millions up to billions of years ago. The most distant supernova observed so far (by Hubble) is called UDD10Wil and it is about 10 billion light-years away. This means its star exploded 10 billion years ago when the universe was only about 4 billion years old. The environment then and the star itself were probably very different from our modern universe today. When the James Webb Space Telescope starts operating in 2018, astronomers will be able to routinely observe the supernovae of the very first stars to form just hundreds of millions of years after the Big Bang.

Supernova shockwaves, although they travel very fast, about 40,000 km/h, they do not approach light-speed. A supernova shockwave was captured during its first minutes for the first time this year. Just before all or most of the star blows apart, a shockwave starting in the center of the star reaches its surface and expands outward into space, accompanied by a flash of light. It works much like a mechanical shockwave such as thunder rumbling through Earth's atmosphere, except that a supernova shockwave travels through the much more disperse and often ionized gases of interstellar space. To compare medium densities, Earth's atmosphere contains about 3 x 1018 molecules per cubic centimeter (cm3), while interstellar gas contains just 1 atom per cm3. The densest nebula contains about 10,000 molecules per cm3. What makes the shockwave visible are gamma photons accompanying the breakout flash, speeding out of the stars collapsing core and out through the star's surface (traveling at light-speed). By the time one of Earth's telescopes "sees" the distant EM flash, very short wavelength (invisible) gamma photons have stretched into intense visible light.

Why and when does a nebula collapse into a star? A gas/dust cloud exists in a delicate state of hydrostatic equilibrium. It is balanced between two opposing forces: Internal (outward) pressure is exerted by the thermal (colliding) motion of the molecules and atoms themselves and by the interactions between (repulsive) magnetic fields. Ionized gases in the cloud are charged objects. As they move about randomly in the cloud they create magnetic fields, which exert outward pressure. Meanwhile, the cloud particles experience the attractive (inward) force of gravity between them. A balance between these forces is reached, the cloud remains stable, until an outside force applied to the cloud disturbs it. Even a relatively mild disturbance from a gravitational collapse into a star nearby can trigger a local collapse. Regions of dust and gas can also simply collapse under their own increasing density. A typical nebular cloud can be a very productive birthing centre. Averaging about 100 light-years across and containing up to 6 million solar masses (M) of matter, a single nebula can give birth to millions of new stars.

A typical nebular cloud consists mostly of hydrogen. Depending on when and where a nebula exists, it will also contain traces of other larger atoms, molecules and different ionized gases as well, blown in from past supernovae. All the material of our solar system originally collapsed from a nebular nursery cloud much like the Eagle Nebula.

The parts of a nebula that are active star nurseries tend to be dense cold darkly opaque molecular clouds. In these regions, hydrogen exists as an H2 molecule rather than as an ionized atomic gas (a proton). Here, thermal (outward) pressure exerted by the cold hydrogen is minimal, which gives gravitational pressure an edge. It makes local collapses into denser regions possible. Gas and dust collapse inward to an increasingly dense local core that eventually evolves into a dense spinning mass called a protostar. It spins because all the angular momenta of the particles are conserved.

A typical protostar forming within a dense molecular cloud is illustrated left. Two bipolar stellar jets of material are likely the result of interactions between powerful magnetic fields wrapping around the forming star. They spin the material round and round and eject it from the star's magnetic poles.

In 2004, NASA directly glimpsed a protostar (V1647) for the first time. Still swathed in its birth blanket of gas and dust, it was actively accreting mass, a process animated in this brief video clip.

The protostar continues to accrete gas and dust from the surrounding cloud for millions of years until it has enough mass to evolve into a pre-main-sequence star. When it reaches this point it has gained its final mass. The star's mass will determine how it will then evolve over millions to trillions of years to come, as a main-sequence star. It will also determine whether or not it will eventually explode as a supernova. The larger the initial mass of dense gas and dust, the larger the star will be. Because stars vary widely in mass, what defines the birth of a main-sequence star is its core temperature rather than its final mass. Fusion is a very temperature-sensitive reaction. When the core is hot enough to trigger the fusion of hydrogen into helium, the star proper is born and it begins to shine brightly. More massive stars fuse hydrogen at a higher rate than smaller stars do. They go through their fuel faster, which means they don't last as long. The most massive stars last for just a few million years while the smallest ones last for trillions of years.

Next, we will start with the life cycles of the least massive stars - brown dwarf stars, red dwarf stars and small to mid-sized stars like our Sun - and see how they end up. Will our Sun explode into a supernova one day?