Cosmic Evolution Project Submenu Links - click to expand or collapse list

                click on underlined bold text to jump to another page

about

rfield

workshop

 

Universe

origin of the Universe

composition

cosmic web

stardust

supernova

dark-matter-day

Milky Way Galaxy

black-hole

primordial overdensities

Big Bang

Expansion

Early cosmic inflation

protons and neutrons

photons and neutrinos

Stars and Galaxies

Solar System

Sun

Jupiter

Venus

Earth

Moon

Ice Giants

Trans-Neptunian Objects

Asteroid Belt

Moons and Rings

Earth and Geobiosphere

Global Catastrophes

Ocean Science Quest

Darwin In the Garden

Origin of Life

Diversity of Life

Complexity of Life

virus

bacteria

archaea

eukaryotes

plants

fire

Organism Life Cycle

Ecosystem Evolution

Ecosystem Life Cycle

Brains and Tools

good-or-bad

Brain Structure

Brain Cell Building Blocks

Brain chemistry and neuroplasticity

Brain Development

Brain Evolution

Brain Emergent Properties

Consciousness

Brain Life Cycle

Tools

Tools to expand sensory powers

Tools to expand physical powers

Tools to expand mental powers

Artificial Intelligence

Artificial Selection

Socio-economic Evolution

 

 

supernova

This webpage is under construction - last updated 2020-1011.

What do supernovas do?

What is the composition, structure, and evolutionary history of massive stars?

Generally speaking, local properties like composition, density, temperature, and flows and interactions of energy determine the evolutionary history of stars. Gravitational, electromagnetic, and nuclear forces control the lives of massive stars from birth to supernova. Quantum effects like tunneling, electron degeneracy, neutronization, neutrino production, and neutron degeneracy influence their lifetimes and fates. Self gravitating spheres of cold dilute gas and dust ignite and sustain thermonuclear reactions for long periods of stability through the balance of pressures that resist the gravitational contraction of their masses.

Stars evolve continuously because nuclear processes continuously change the composition of stars as well as their thermal properties. One consequence of nuclear reactions is that the mere fact that the number of particles with kinetic energy changes affects the ability to resist contraction and therefore the stability, rates of subsequent reactions, and longevity of stars.

A massive star is born

the core of a main sequence massive star is far less dense than the sun's

A star is a self-gravitating sphere of gas that contracts as it radiates some of its energy into space. The rest of the energy is trapped by the opacity of a star so that the core grows denser and hotter as gravitational potential energy is transformed into thermal energy. By replacing the lost energy that flows to the surface by convection or as radiant energy, fusion maintains the hydrostatic pressure that resists the gravitational contraction of a star.

The sun cannot ignite fusion until the density is high enough for a few protons to overcome the repulsive Coulomb force of like charges. A massive star achieves core stability by igniting the CNO cycle before the core density is very high. The secret is that the CNO cycle is far more temperature dependent than the proton-proton reaction. A massive star reaches a higher core temperature than the sun before it contracts enough to produce a high core density.

What does the sun do?

The sun was born with a hot dense core ideal for thermonuclear fusion as a result of the energy and pressure provided by the gravitational collapse of the massive cool dilute gas and dust that formed it. The sun will live about ten billion years on the main sequence, fusing hydrogen into helium in a process that begins with the rare fusing of two protons that did not have enough kinetic energy to overcome their electrical repulsion, but were helped by quantum tunneling.

The sun is stable throughout this time with its core slowly growing hotter and denser and its surface growing larger, hotter, and brighter. The proton-proton reaction produces neutrinos that quickly escape unimpeded, but only a small fraction of the energy is lost this way. If most of the energy were lost by neutrinos, the sun would contract quickly and live a brief intense life.

In the core of the sun, most of the energy generated by fusion is released as photons that are quickly absorbed in the opaque depths of the fully ionized plasma of nuclei and elections. This provides the energy to sustain the pressure necessary to oppose the force of gravitational attraction that would follow the loss of energy radiated from the surface. This energy takes millions of years to flow through the radiative zone. The outer part of the sun is so opaque that radiant energy is trapped and convective heat flows take over quickly transporting the energy to the surface where it is radiated into space.

When the sun has exhausted its core hydrogen billions of years from now, it will become a red giant and burn helium, but very little of the few elements like carbon and oxygen that it synthesizes in its post main sequence life will escape to enrich interstellar medium for the benefit of future generations of stars.

Main sequence sun-like stars synthesize helium in three typical steps consuming four protons, two of which turn into neutrons and positrons. This reaction is typical in smaller stars like the Sun. It is so difficult to combine protons that repel each other electrically that the main sequence lasts for millions or billions of years until most of the core protons is able to join another one in the first step of this process.

Positrons are the positively charged anti-particle of electrons. Positrons and electrons annihilate each other when they meet and release radiant energy. The energy released through gamma ray production and through positron-electron annihilation keeps the star from cooling and contracting much as it radiates starlight to space.

Step one is two hydrogen nuclei or protons combine and emit a positron producing an isotope of hydrogen whose nucleus is called a deuteron containing a proton and a neutron. This is the slowest step. When a third proton joins the deuteron, helium-3 is produced, a new element with two protons, but a less common isotope of helium than the common one with two neutrons. This process happens again and the two helium-3 nuclei can fuse forming the highly abundant helium-4 isotope, producing another positron and freeing two protons. So six protons got together and two came out, so four protons turned into one helium-4.

What do massive stars do?

Massive stars behave very differently than the sun. Stars with 25 solar masses end with violent supernova explosions, enriching interstellar medium with a wide variety of elements synthesized during their short intense life, leaving a neutron star remnant or a black hole where its high density core collapsed. It turns out that neutrinos and neutrons play major roles in the life and death of massive stars.

Having 25 times more hydrogen to burn than the sun does not give them longer lives. Instead it produces a higher core temperature that triggers a highly temperature dependent process based on a small amount of carbon, nitrogen, and oxygen in the star. CNO cycles are much more efficient at capturing protons and emitting alpha particles which are helium nuclei. Again a small amount of energy is lost to neutrinos.

Back to top

CNO process synthesizes helium from hydrogen on the main sequence of a blue star

Massive main sequence stars also make helium in a series of steps that also consume four protons, two of which turn into neutrons and positrons. Positrons are also produced and release radiant energy when they encounter an electron. Unlike smaller stars like the Sun, massive stars have various CNO cycles. One example is  the cyclical steps of isotopes of 12C 13N 13C 14N 15O 15N and back to 12C by ejecting one 4He.

The CNO process is so efficient that it stabilizes the contraction of the gases while the core density is much less than that of the sun. This is possible because the moderately higher temperature of the greater contracting mass triggers the CNO cycles. Because of the greater radius of the star and the greater production of energy, the photons are unable to transport the energy efficiently from the core and instead of a radiative core and convective envelope like the sun, massive stars have convective cores and radiative envelopes.

Back to top

despite all their fuel, massive stars live short lives compared to the sun

There is a rule of thumb that main sequence luminosity varies as mass to the 3.5 power. A 25 solar mass star may be 78,000 times brighter than the Sun. Since it only has 25 times as much hydrogen as the sun, its life on the main sequence may be thousands of times briefer – six million years rather than ten billion years.

The following semilog bar chart shows the incremental time (delta time) preceding each of the 13 snapshots of stellar evolution that were modeled in a 25 solar mass pre-supernova simulation data provided to us by UCSC Professor Stan Woosley. The data is his, the interpretation is ours.

 

Back to top

massive stars contract quickly when core hydrogen is depleted

Like the sun, the depletion of hydrogen from the core removes the energy source that supports the pressure that resists gravitational contraction. Helium burning is ignited after the core contraction has increased the central density by a factor of 90 and the conversion of gravitational potential energy to thermal energy has increased the central temperature by a factor of three.

The following two tables show changes in mass, surface radius, average density, effective temperature, and peak wavelength for the 13 snapshots of evolutionary history in the simulation data provided to us. Some of the properties are expressed as fractions or multiples of the 4.6 billion year old sun. The values of the current age sun are indicated below the tables. A set of four charts plots the values in the tables on linear or logarithmic scales as needed.

The discussion and conclusions described throughout this webpage are based primarily on our understanding and interpretation of simulation data.

 

A red supergiant grows up

post main sequence massive stars evolve into red supergiants

The end of hydrogen burning leads to a rapid increase in the central density and temperature. This triggers helium burning which stabilizes core contraction. So much energy is generated that the opacity of the envelope of the star causes it to heat up and expand rapidly. The surface temperature drops when the star becomes a red supergiant star with a volume a million times greater than it had on the main sequence as a blue star.

Because the surface area is about ten thousand times greater than the main sequence surface area, the supergiant star can radiate the energy into space so efficiently that the surface temperature drops by a factor of ten and its blackbody peak wavelength increases by a factor of ten, changing its appearance from blue (actually peaked in the ultraviolet) to red (actually peaked in the very near infrared).

When an element is depleted in the center, the central density and temperature rises igniting some of the fusion products and elevating the surface luminosity. Even though the surface luminosity increases, the surface temperature decreases because the energy is distributed over a giant surface area as the energy heats and expands the envelope.

The helium burning phase is shorter than the main sequence phase – less than a million years under these conditions and the surrounding gas is so hot and dense that hydrogen burning takes place in a shell around the core. The primary reaction during helium burning is the triple alpha process in which carbon is synthesized from three helium nuclei. A fourth alpha particle can be captured to synthesize oxygen.

When the core helium is depleted, an enormous amount of carbon and oxygen has been synthesized while the amount hydrogen has decreased significantly. The energy released is photons rather than neutrinos and heats the entire interior, turning the star into a red giant that grows to 100 times its original radius and a million time lower average density. Its radius is comparable to the radius of Jupiter’s orbit.

As successive elements form, they fuse until being depleted at the core. The primary sequence of elements burning at the center are hydrogen, helium, carbon, neon, oxygen, and silicon. As an element is depleted at the center, it may begin to burn in a shell surrounded the original core, so that over time, many shells are burning.

This process continues until iron group elements form. Since they cannot burn because of a property of element binding energy, fusion stops at the core leading to catastrophic core collapse due to gravitational pressure that is no longer balanced by the generation of heat.

The two tables below are additional columns of the two previous tables that are based on the 13 snapshots of the evolutionary history of a 25 solar mass pre-supernova star. These tables show properties of the star including central density and temperature, effective luminosity, and neutrino luminosity as well as comparisons to the current age sun. A set of four charts plots the values in the tables on linear or logarithmic scales as needed.

 

triple alpha process synthesizes carbon from helium as the star evolves into a red giant

The triple alpha process synthesizes carbon from three alpha particles. Alphas are the highly abundant helium nuclei synthesized minutes after the Big Bang and by main sequence stars. Note that two helium nuclei cannot form a stable form of beryllium. If the beryllium does not acquire a third alpha particle quickly, it decays back to the two helium nuclei. A stable isotope of beryllium is produced by a different process.

The afterlife of a massive star is very different from the fate that the Sun will experience five billion years from now when it exhausts its core hydrogen and leaves the main sequence. Massive stars go on to synthesize most of the most abundant elements and disperse them in a supernova explosion as they collapse into neutron stars or black holes.

Stars like the Sun synthesize carbon and oxygen from helium and become red giants rather than red supergiants. Their central density is high, but their average density is ten million times less dense than the Sun. When core fusion ends, they do not disperse elements in massive explosions. Instead, they condense into white dwarf stars one million times denser than the Sun.

The contraction of less massive stars is ultimately halted by electron degeneracy pressure. The central density of a carbon-oxygen degenerate dwarf star is somewhat higher and the star does not do any fusion, so it slowly cools for billions of years until it dims and evolves into a black dwarf, a fate that has not been reached by any stars yet.

 

 

red supergiants synthesize every element from carbon (z=6) to zirconium (z=40)

What elements do massive stars form? Massive stars make everything else from carbon (z=6) to zirconium (z=40). Some of the most abundant are the building blocks of life: carbon, nitrogen, oxygen, sulfur, and phosphorus. Some are common in the oceans where life evolved: sodium and chlorine.

The sun is over 98% hydrogen and helium. Carbon, nitrogen, oxygen, and magnesium comprise 98% of the rest, the so-called metals.

The solid Earth consists mostly of ten elements. In fact over 90% of the Earth's mass is oxygen, iron, silicon, and magnesium. Add nickel, aluminum, sulfur, and chromium and you have accounted for over 99% of the Earth's mass.

The entire process of short lived massive stars producing and dispersing metals that enrich the interstellar medium early enough to be incorporated in later generations of long lived stable smaller stars is critical to the formation of Earth-like planets with enough time, material, and solar energy for complex intelligent life to evolve.

 

computer simulations model the evolutionary history of stars

UCSC Professor Stan Woosley has studied the lives of massive stars that die in Type II supernova explosions for decades. He was generous enough to provide us with massive amounts of data on 25 solar mass stars from his simulation code as well as guidance and advice. We analyzed data from 13 snapshots or time slices from ZAMS to pre-supernova that included radially varying properties like total mass, density, temperature, luminosity, and element composition and the associated nuclear reactions. Neutrino luminosity vs. time is very interesting.

Our studies revealed patterns, many more complex than we have had time to explore this summer. The underlying processes are numerous and this report attempts to describe a few key features. This work is a continuation of work done to produce lthe November 2019 We are stardust talk. Our goal was to understand the processes that produce the naturally occurring elements of the Periodic Table. We focused on massive stars this time. These stars produce significant amounts of the elements that are most common and/or important to the formation of long lived smaller stars as well as rocky planets and life.

We learned a lot about the critical role of photons, neutrons, neutrinos, and quantum effects throughout the lives of stars. Another big idea is that the lives of stars are ruled by the breakup of particles as well as their fusion.

Let’s consider the roles of neutrinos, dark matter, and quantum effects in the universe. Without the critical role of neutrinos in the asymmetry of matter and antimatter, matter would not exist. Neutrinos interact with dark matter which preserved quantum density fluctuations that eventually formed stars and galaxies. The fact that main sequence sun-like stars do not lose much energy to neutrinos is critical to small stars living long enough for life to evolve. Quantum tunneling plays a key role in the main sequence lives of the small stars that burn hydrogen primarily through the proton-proton reaction.

The fact that the CNO cycle is so sensitive to temperature enables massive stars to burn efficiently on the main sequence and live short lives. The energy loss through neutrino production from carbon burning on insures that stars go through the post-main sequence lives quickly. The high neutrino flux from neutron star formation is key to the explosion of supernovas.

The set of four charts below show the evolutionary history of the enclosed mass and density as functions of the logarithm of a star’s radius based on a 25 solar mass pre-supernova simulation model provided by UCSC Professor Stan Woosley. More of the simulation results will be presented and discussed throughout this webpage.

ZAMS refers to the zero age main sequence state of a star and hedep refers to the point at which helium burning in the core stops because helium is depleted. This sets the stage for the subsequent stages of core burning, beginning with carbon ignition cign and ending with the pre-supernova state presn. The charts show that in the later stages of evolution, mass grows faster with radius because the density of the core rises. The charts also show that the surface radius increases until helium depletion occurs.

The set of four charts below show the evolutionary history of the temperature and luminosity as functions of the logarithm of a star’s radius based on the 25 solar mass pre-supernova simulation. The charts show that the core temperature and luminosity increase throughout the star’s evolution. The charts also show that the surface temperature increases and the luminosity decreases until helium depletion occurs.

 

Neutrinos earn their metals

neutrinos rapidly dissipate energy from massive stars late in their lives

Let’s consider the roles of neutrinos, dark matter, and quantum effects in the universe. Without the critical role of neutrinos in the asymmetry of matter and antimatter, matter would not exist. Neutrinos interact with dark matter which preserved quantum density fluctuations that eventually formed stars and galaxies. The fact that main sequence sun-like stars do not lose much energy to neutrinos is critical to small stars living long enough for life to evolve. Quantum tunneling plays a key role in the main sequence lives of the small stars that burn hydrogen primarily through the proton-proton reaction.

The fact that the CNO cycle is so sensitive to temperature enables massive stars to burn efficiently on the main sequence and live short lives. The energy loss through neutrino production from carbon burning on insures that stars go through the post-main sequence lives quickly. The high neutrino flux from neutron star formation is key to the explosion of supernovas.

The post main sequence lives of massive stars would be very different if most of the energy had to be radiated away by photons. Neutrinos can escape the interior of a star at the speed of light because they do not interact electromagnetically with the positive nuclei and negative electrons and their rest mass is negligible, so gravity does not trap them. About 15% of the energy lost to space when carbon burning ignites is via neutrinos. By the time the star becomes a pre-supernova, neutrinos transport ten million times more energy to space than photons do, a rate that has increased by a factor of a million in the final year of the star’s life.

carbon burning and the building blocks of life and planets

The carbon in the core burns when the helium is depleted and the core has contracted further to another seven fold increase in density (12,000 times water) and twice the temperature (500 million kelvin). The carbon core is enclosed in shells of burning helium and hydrogen. While the carbon burning does not produce neutrinos, 15% of the energy does escape as neutrinos due indirectly to pair production.

At high temperatures, a pair production process may turn a high energy gamma ray into an electron and positron. In 1 out of 1019 pair productions, a weak interaction replaces them with a neutrino / anti-neutrino pair that escape the star with energy loss comparable to the output from carbon fusion. Neutrino losses play an increasingly important role over time by forcing a star to burn its fuel at ever higher temperatures to maintain internal pressure.

Let’s explore the typical nuclear processes in the concentric shells of the massive star after it leaves the main sequence. We already talked about the triple alpha process forming carbon. At high densities and temperatures, two carbon nuclei can fuse to form a variety of common and important elements including neon, sodium, oxygen, and magnesium. The simple example illustrated here is the formation of magnesium with the release of energy. In this case all twelve of the protons and all twelve of the neutrons present in the original two carbon nuclei remain in the magnesium nucleus.

The neon, oxygen, and silicon burning stages progress as the central density and temperature rise and as a result of composition changes which deplete the supply of fuel available to sustain hydrostatic pressure.

neon and oxygen burning synthesize silicon, phosphorus, sulfur, and magnesium

When a neon burning shell forms, oxygen and magnesium are often created. One process produces an alpha particle (helium nucleus) and the other consumes one.

Similar to the carbon burning, two oxygen nuclei can fuse together to produce new elements. Oxygen burning synthesizes silicon, phosphorus, sulfur, and magnesium. The three most common processes are shown here with their relative rates in percent. As a bonus, these reactions produce other particles like alphas, protons, and neutrons that may participate in other reactions. Remember that a lot of processes are needed to produce the very large number of elements in the Periodic Table.

 

Back to top

silicon burning uses the alpha process to synthesize elements up to zinc (z=30)

Earlier we saw a list of alpha ladder process that ended in iron. This diagram shows that silicon burning uses alphas to produce elements from sulfur to zinc. Each step increases the number of protons by two and the number of neutrons by two. So every other element in this section of the Periodic Table can be made by this chain of nuclear events.

The binding energy per nucleon has a broad peak around iron group elements. Binding energy is the energy released when particles are bound together. For nuclear reactions, the energy comes from a  mass reduction as in E = Mc2. When the curve slopes downward, the total binding energy may be less than the energy of the particles that were combined to form them. An explosion may be the only way to provide the additional energy to produce these heavier elements.

 

Back to top

A neutron star is born

photodisintegration and electron capture turn collapsing core into neutron star

Photodisintegration breaks iron-56 into 13 alpha particles and four neutrons. Electron capture turns an alpha particle into four neutrons. Loss of electron degeneracy pressure causes core to collapse. Neutron degeneracy pressure halts contraction producing a stable neutron star.

Photodisintegration breaks iron-56 into 13 alpha particles and four neutrons. Electron capture turns an alpha particle into four neutrons. Also called neutronization, the process also releases neutrinos. Loss of electron degeneracy pressure causes core to collapse. Neutron degeneracy pressure halts contraction producing a stable neutron star.

 The radius of the core shrinks from about 1000 km to 10 km in a matter of seconds. The star converts enough gravitational energy to radiant energy for the star to be millions or billions of times brighter than the Sun for about one year.

The neutron star has a thin outer crust of ions and electrons. The inner crust consists of electrons, neutrons, and nuclei. The outer core contains a neutron-proton Fermi liquid and an electron Fermi gas. It is believed that the inner core is dense enough for baryons to decompose into a quark gluon plasma.

 

Back to top

rapid neutron capture synthesizes elements up to zirconium (z=40)

In a fraction of a second, the absorption of neutrons from the flux from the collapsing core forms elements beyond zinc. The capture of a large number of neutrons by a nucleus produces an extremely unstable isotope. Beta decay reduces the number of neutrons and increases the number of protons, resulting in a new element with the same number of nucleons as the nucleus had after neutron capture.

 

Back to top

neutron star mergers synthesize elements from niobium (z=41) to uranium (z=92)

If neutron star has a binary partner from another supernova and they merge, they will form and disperse all of the naturally occurring elements beyond zirconium, that is from niobium (z=41) to uranium (z=92).

The so-called r-process or rapid neutron capture process occurs in the debris ejected in a neutron star merger. Many neutrons are freed and quickly increase the size of nearby nuclei. These nuclei are overloaded with neutrons compared to protons and are unstable. When a neutron decays, a proton and an electron are produced. The escaping electron is a beta particle. The atomic number Z increases by one, creating a nucleus of a new heavier element. This process can be repeated many times in a second, forming many new elements.

This table is a placeholder for a table showing some of the important elements present at each evolutionary stage of the simulation model. Because of various normalization factors related to the way the simulation is run, this table is fundamentally inaccurate and needs to be corrected to match the actual predictions of the simulation.

These two charts show the relative amounts of the main core elements vs. radius at the end of the main sequence.

 

Back to top

what if a supernova remnant is a black hole rather than a neutron star?

The supernova explosion occurs as the core of the star is collapsing. If the neutron star is more than about three solar masses, then the neutron degeneracy pressure is not sufficient to halt the collapse and the star continues to collapse to form a black hole. Since the star forms a neutron star before it degenerates into a black hole, the supernova event is generally the same regardless of whether the remnant is a neutron star or a black hole.

 

Back to top

Neutrinos have a blast and we inherit the wind

neutrinos deposit energy in the shell surrounding the collapsing core

The flux of neutrinos from the core is so intense that enough neutrinos are absorbed to propel the supernova explosion. At the pre-supernova stage, the neutrino luminosity is 3 x 1015 the luminosity of the sun whereas the optical luminosity of the star has not changed much since it reached its full size as a red supergiant shortly after helium depletion. Throughout its post-main sequence life, the star’s luminosity of about 250,000 solar luminosities has not even doubled from its value at the end of its main sequence life as a blue giant.

 

Back to top

additional elements are synthesized during a supernova explosion

Short-lived radioactive isotopes that form during the explosion may decay on relatively short time scales to form new elements.

 

Back to top

additional elements are synthesized in the nebula formed from the explosion

Short-lived radioactive isotopes that form during the explosion may decay relatively quickly to form new elements after the envelope has been propelled into space.

 

Back to top

the interstellar medium is enriched with the building blocks of planets and life

The debris from the explosion can enrich the interstellar medium where the next generation of stars is formed along with their planets and possibly even life. Retention of the debris is enhanced by the strong gravitational attraction in a massive galaxy, most of whose mass is dark matter.

 

Back to top

additional information

animations and illustrations

new
 

simulations and models

new

 

Wikipedia articles

new
 

more resources

more videos

new

 

 

 

Related Content