Core-Collapse Supernovae and the Formation of Heavy Elements

NASA and the European Space Agency - via Wikimedia Commons.

The main elements in the human body are carbon, hydrogen, nitrogen, oxygen, phosphorus, and sulfur. The acronym, CHNOPS stands for these six elements, each of which is uniquely essential to life, yet the most common among them are four specific elements: carbon, hydrogen, nitrogen, and oxygen. Hydrogen has one proton, carbon six, nitrogen seven, and oxygen eight. Without any one of these elements, life on Earth would be dramatically altered or even nonexistent. 

Most stars originally fuse hydrogen into helium,* leaving two abundant, obviously existing elements in the universe. But what about nitrogen, oxygen and carbon? How is it possible that, while the main sequence stars are fusing hydrogen into helium, we have carbon, nitrogen, oxygen, phosphorus, and sulfur? There is more to the universe than hydrogen fusion; welcome to the post-main sequence** life cycle of a star, and the core collapse of large stars.

The Life Cycle of a Star

The size of a star determines its life cycle – how long it will last and how it will die. The life cycle of any star first begins with the gravitational collapse of a molecular cloud, a cloud of predominantly hydrogen gas that, once provoked by a gravitational imbalance – such as a passage of a nearby star – begins to condense, producing a positive feedback loop that causes a small density of hydrogen, which has a slightly greater gravitational force, to gradually accrete more hydrogen from the surrounding molecular cloud. The collapse of the molecular cloud focuses on a particular point of greater density, soon forming a hot, gaseous body – a protostar. Protostars do not have the nuclear capabilities of a main sequence star, and they are considered “protostars” because they are still accreting mass from their parent molecular clouds. A typical low-mass star is a protostar for around 500,000 years.

As the protostar, having gathered much of the available gas in the molecular cloud, nears its final mass, it becomes a pre-main sequence star – a star that, although yet to achieve nucleosynthesis, has grown to near its final mass. Some such stars may ever reach nucleosynthesis, becoming brown dwarfs – failed stars that can never perform full nuclear fusion to helium, but instead from hydrogen to deuterium, a “heavy hydrogen” atom with one proton and one neutron. Some pre-main sequence stars do not even achieve hydrogen to deuterium fusion, becoming sub-brown dwarfs. One such hypothesized sub-brown dwarf (or failed brown dwarf) is Jupiter, but most scientists do not support this position.

Once the core in a star reaches 10 million degrees kelvin, it will achieve nucleosynthesis and begin the process of fusing hydrogen to helium. At this point, such a star enters the main sequence.

Once a star is in the main sequence, the radiation pressure from the fusion in its core balances out the gravitational pressure attempting to collapse the star, leading to hydrostatic equilibrium and keeping the star from either imploding or exploding. It is in this stage that a star tends to be the most stable. 

Any star will fuse hydrogen to helium for a significant portion of its lifespan, generally until it runs out of hydrogen in its core. Once a star ceases fusing hydrogen to helium in its core, it will leave the main sequence.

After a star’s time in the main sequence, the rest of its life depends entirely upon its mass. Red dwarfs to sunlike stars, which have a mass anywhere from .08 to 8 solar masses, will transition in a similar way the sun will. Larger stars – those larger than eight solar masses – will evolve much differently.

Stars smaller than eight solar masses

Stars of this size – smaller than eight solar masses – may encompass 99.99957%*** of the stars in the visible universe, and generally exhibit sunlike features – at varying degrees. The larger of these stars will produce brilliant planetary nebulae, and the smallest may survive in the main sequence for trillions of years, whereas the largest will for less than the sun will.

As the star begins to run out of hydrogen to fuse in its core, it begins to change. For example, the sun is projected to grow 1% brighter every 100 million years for the next five billion years. And once the sun runs out of hydrogen in its core, it will, for a short time, not perform nuclear fusion, thereby making the forces of radiation and gravity unbalanced; this results in an unhindered gravitational force contracting the star’s mass, the contraction of which heats up both the core and the hydrogen region surrounding the core, leading the area surrounding the core to begin performing hydrogen fusion again, and eventually, the now much smaller core of the star to begin performing helium fusion. This stage will continue until the core runs out of helium, which will result in greater contraction of the core as the forces between gravity and radiation become unbalanced once again.

For stars under 8 solar masses, carbon fusion cannot be achieved, and thus the star will continue to contract until it becomes a white dwarf, the remnant core of the star. The rest of the star will be ejected from the white dwarf, forming a planetary nebula, a rich and generally spherical, elliptical, or bipolar region of gas that generally lasts for a few ten thousand years.

Stars larger than eight solar masses

The most significant distinction between stars of this size – eight or greater solar masses – and the smaller stars is that these larger stars can fuse beyond hydrogen and helium. Such stars also exist for much less time than smaller stars: the largest stars, for example, exist for only a few million years, while the smallest main sequence stars possible – red dwarfs – can perform fusion for upwards of six to twelve trillion years.

As stars greater than eight times as massive as the sun pass helium fusion, they continue to fuse all the way to iron in this order: carbon and oxygen fusion; then magnesium, oxygen, and neon fusion; then silicon and sulfur fusion, which fuses to iron – an element that is incapable of fusion in such an environment. The fusion to iron requires more energy than it produces, thereby dooming this large star and leading it to among the most destructive ends in the universe: supernova.

The End?

Once this star fuses to iron and produces iron ash, not enough energy is immediately available to fuse iron. As a result, some matter flows towards the core of the star as it begins to collapse.

As fusion has since ceased in the core of the star, the forces between radiation and gravity become, once again, unbalanced. Once the core becomes entirely iron and reaches critical mass, it implodes. The implosion of such a large star is profound, so much so that it rips the star apart, heating the plasma and speeding it up to several percent the speed of light. High energy radio and UV rays are also released immediately after the explosion and can affect ozone layers for hundreds of light years. 

The luminosity of a supernova immediately after core collapse can be comparable to that of an entire galaxy, and nearby supernovae could even be as bright as the full moon. When, for example, a star named Betelgeuse, 642 light years away, goes supernova, it will shine as bright as the half moon and will be easily visible during the day for weeks after the blast. The remnants of supernovae will remain visible long thereafter, as the shockwave of the supernova – known as supernova remnants – collects and excites gas in the interstellar medium. The Crab Nebula, for example, is a remnant of a supernova that occurred in 1054.

Processes behind trans-iron elements

As we know that even the largest stars fuse only to iron, we now ask the question: how did the trans-iron heavy elements, like plutonium or uranium, form? How could we have atoms with 118 protons when the largest and stars directly produce atoms with only 26? 

There are three processes theorized to be responsible for trans-iron elements: the r-process, the p-process and the s-process. 

The r-process – short for the rapid neutron capture process – is primarily responsible for around half the atomic nuclei with more than 26 protons. It is characterized by a rapid succession of neutron captures – essentially, atomic nuclei capture a ton of neutrons in a short period. The seed nuclei – a nucleus that serves as a starting point for fusion reactions to occur – is generally iron. The rapid neutron captures must be rapid enough so that 𝛃- decay – the conversion of a neutron to a proton through extreme processes – occurs at a slower rate than a new neutron is captured. The nucleus becomes increasingly neutron-rich as it remains in the supernova’s extreme environment; as it reaches its peak neutron count, it will gradually undergo 𝛃- decay, resulting in a higher atomic number but a near-consistent atomic mass, as all that is lost in negative beta decay are an electron and an antineutrino, both of which have very little mass compared to the neutron or the proton.

The p-process – short for the proton process – results in neutron-deficient isotopes of elements 34-80 (selenium through mercury). During the supernova, protons are added to an atomic nucleus, again generally that of iron, in a process known as a proton capture reaction. Like the r-process, the p-process requires significant temperature and pressure, both of which are present in supernovae.

The s-process – short for the slow neutron capture process – is also responsible for almost half of the atomic nuclei heavier than iron. The capture is similar to the r-process, but the s-process occurs much more slowly than the r-process – slow enough that radioactive decay (whether alpha or beta decay) actually can occur before another neutron is captured, thereby limiting how large the nucleus can get. The s-process often occurs in supernovae of the largest of the asymptotic giant branch stars – red giants and supergiants with solar masses of .2-10. Because this process is slower moving than the r-process, the s-process cannot produce ultra-heavy elements like uranium and thorium.

Wrapping it up

Supernovae are among the most explosive events in the universe. Let us rejoice in the fact that our sun will never go supernova and is not volatile enough to become a red giant or an unstable star anytime soon. As always, take care and stay curious, everyone.

* Although most stars fuse hydrogen to helium, some bodies still qualify as stars – brown dwarfs, for example – do not fuse hydrogen at all, or fuse it into deuterium, or heavy hydrogen – a hydrogen atom with two neutrons.

** The main sequence encompasses the vast majority of stars in the present universe; stars in this portion of the Hertzsprung-Russell diagram are in the process of fusing hydrogen in their core to helium. The star is one such example of a main sequence star.

*** Obtained through calculations of the stellar mass function; these percentages are still the feature of great debate in the scientific community, and are mathematical estimates, not observed values. 

If you have any questions, comments, or corrections, please comment on this post or email with your concerns. Thank you.


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