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Stellar Evolution: Stages in the Life Cycle of Stars

Dr. Thomas Swan is a published physicist who received his PhD in nuclear astrophysics from the University of Surrey.

The physical characteristics of stars are usually quoted relative to our Sun (pictured).

The physical characteristics of stars are usually quoted relative to our Sun (pictured).

The Physical Features of Stars

Before learning about the life cycle of stars and the astrophysical processes that are involved, it is useful to know something about the fundamental features of all stars in the universe.

Stars are luminous spheres of burning gas that are between 13 and 180,000 times the diameter (width) of the Earth. The Sun is the nearest star to Earth and is 109 times its diameter. For an object to qualify as a star, it must be large enough for gravitational forces to have created enough pressure in its core to trigger nuclear fusion.

The surface temperature of the Sun is 5,500 °C, with a core temperature as high as 15 million °C. For other stars, the surface temperature can range from 3,000 to 50,000 °C. Stars are predominantly composed of hydrogen (71%) and helium (27%) gases, with traces of heavier elements such as oxygen, carbon, neon and iron.

Some smaller stars have lived since the earliest era of the universe, showing no signs of dying after more than 13 billion years of existence. Others (the biggest and hottest) live only a few million years before burning through all their fuel. These stars can have 300 times the mass (weight) of the Sun and be 9 million times as luminous. Conversely, the smallest stars can be 1/10th of the mass and have 1/10,000th of the luminosity.

Without stars we would simply not exist. These cosmic behemoths use nuclear fusion to produce the elements that planetary systems and living beings are made of (carbon, oxygen, iron, etc.). The next sections will describe this process by outlining the different stages in the life cycles of stars.

A region of the Carina Nebula called Mystic Mountain in which stars are forming.

A region of the Carina Nebula called Mystic Mountain in which stars are forming.

The Birth of Stars

Stars are born when nebulous clouds of hydrogen and helium gas coalesce under the force of gravity (see above). Sometimes a shock wave from a nearby supernova is required to produce areas of high density in the cloud.

These dense pockets of gas within the cloud contract further under gravity while accumulating more material. The contraction heats up the material, causing an outward pressure that slows the rate of gravitational contraction. This state of balance is called "hydrostatic equilibrium" and a dense pocket of gas that is in this state is called a "protostar."

Contraction comes to a complete stop once the core of the protostar becomes hot enough for hydrogen to fuse together to make helium in a process called nuclear fusion. At this point, the protostar becomes a "main sequence star."

Star formation often occurs in gaseous nebulae in which the density of the nebula is great enough for hydrogen atoms to chemically bond to form molecular hydrogen. These nebulae (or clouds) are often called stellar nurseries because they contain enough material to produce several million stars, leading to the formation of star clusters (see below).

A star cluster in the Carina Nebula, which is 8,500 light years from Earth.

A star cluster in the Carina Nebula, which is 8,500 light years from Earth.

The Life of Stars

After a star has formed and is on the "main sequence" in its evolution, it will be fusing hydrogen in its core. Hydrogen is the simplest form of atom, with one positively charged particle called a proton orbited by a negatively charged electron, although the electron is lost due to the intense heat of the star. This means that the hydrogen is "ionized."

The stellar furnace causes the ionized hydrogen (i.e., protons or 1H) to slam into each other. At core temperatures above 4 million °C, they fuse together to form helium (4He), releasing their stored energy in a process called nuclear fusion (see below).

During fusion, some of the protons are converted into neutral particles called neutrons in a process called radioactive decay (beta decay), which emits a small positron particle to carry away the proton's positive charge. The energy released in fusion heats the star further, causing more protons to fuse in a self-sustaining reaction.

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The fusion reaction fueling the universe, in which four hydrogen nuclei (protons) fuse into one helium nucleus (He).

The fusion reaction fueling the universe, in which four hydrogen nuclei (protons) fuse into one helium nucleus (He).

Nuclear fusion continues in this sustainable way for between a few million and several billion years, until the star runs out of fuel (sometimes longer than the age of the universe, which is 13.8 billion years).

Contrary to expectations, the smallest stars (red dwarfs, see below) live the longest. Despite having more hydrogen fuel, large stars (giants, supergiants, and hypergiants) burn through it quicker because the stellar core is hotter and under greater pressure from the weight of its outer layers. Smaller stars also make more efficient use of their fuel because it is circulated throughout the star's volume via convective heat transport.

Binary red dwarf stars (Gliese 623) that are 26 light years from Earth. The smaller star is only 8% of the Sun's diameter.

Binary red dwarf stars (Gliese 623) that are 26 light years from Earth. The smaller star is only 8% of the Sun's diameter.

If the star is large enough and hot enough (core temperature above 15 million °C), the helium produced in nuclear fusion reactions will also be fused together to form heavier elements such as carbon, oxygen, neon, and finally iron.

Elements heavier than iron, such as lead, gold, and uranium may be formed by the rapid absorption of neutrons, which then beta decay into protons. This is called the r-process for "rapid neutron capture," which is believed to occur in supernovae.

The Death of Stars

Stars eventually run out of fuel. This first occurs in the stellar core as this is the hottest and heaviest region. Without the heat generated by fusion, the core begins a gravitational collapse, which consequently creates more heat and pressure, triggering fusion in the outer layers of the star where hydrogen fuel still remains. As a result, these outer layers expand to dissipate the heat being generated, becoming large and highly luminous. This is called the "red giant" phase of a star's evolution (see below). Stars smaller than about 0.5 solar masses skip the red giant phase because they cannot become hot enough.

Betelgeuse is a red supergiant that is a thousand times larger than the Sun.

Betelgeuse is a red supergiant that is a thousand times larger than the Sun.

The contraction of the stellar core eventually becomes so extreme that the outer layers of the star are thrown off altogether, forming a planetary nebula (see below; although there are no actual planets). The core stops contracting once the density reaches a point where electrons are prevented from moving any closer together. This is called "electron degeneracy" and it is determined by a physical law called Pauli's Exclusion Principle. The exposed core is called a white dwarf, which gradually cools to become a black dwarf.

A planetary nebula (the Helix Nebula) expelled by a dying star.

A planetary nebula (the Helix Nebula) expelled by a dying star.

Stars of more than 10 solar masses will typically undergo a more violent expulsion of the outer layers called a supernova (see below). In these larger stars, the gravitational collapse will be such that greater densities are reached within the core. Densities high enough for protons and electrons to fuse together to form neutrons may be reached, releasing the energy sufficient for supernovae. The super-dense neutron core left behind is called a neutron star, which is held up by neutron degeneracy in the same way white dwarfs are sustained by electron degeneracy.

A supernova remnant (Crab Nebula).

A supernova remnant (Crab Nebula).

Massive stars above about 20 solar masses will leave behind cores/remnants that are too dense and heavy for neutron degeneracy to prevent them from collapsing further into a gravitational singularity. These stars end their lives as black holes.

The expulsion of a star's matter returns it to the cosmos, providing fuel for the creation of new stars. As larger stars contain heavier elements (e.g. carbon, oxygen, and iron), supernovae seed the universe with the building blocks for Earth-like planets and for living beings such as ourselves.

Tracking the Evolution of Stars

As stars progress through life, their size, luminosity and radial temperature change according to predictable natural processes that can be plotted on Hertzsprung-Russell diagrams. This section will describe that evolution.

How a Protostar Becomes a Star

Prior to igniting fusion and becoming a main sequence star, a contracting protostar will reach hydrostatic equilibrium at around 3,500 °C. This particularly luminous state is proceeded by evolutionary stages called the "Hayashi track" and the "Henyey track," which can be seen on the diagram below.

A Hertzsprung Russell diagram that tracks the early evolution of the Sun from protostar to main sequence star. The evolution of heavier and lighter stars are compared.

A Hertzsprung Russell diagram that tracks the early evolution of the Sun from protostar to main sequence star. The evolution of heavier and lighter stars are compared.

On the Hayashi track, the protostar is still gaining mass from the stellar nursary, and the accumulation of this material is increasing its opacity, which prevents the escape of heat via light emission (radiation). Without such emission, the protostar's luminosity begins to decrease (see diagram above). However, this cooling of the outer layers causes a steady contraction that puts pressure on the core, heating it up. To efficiently transfer this heat, the protostar becomes convective, meaning that hotter material moves toward the surface.

If the protostar has accrued less than 0.5 solar masses, it will remain convective and will stay on the Hayashi track for up to 100 million years before igniting hydrogen fusion and becoming a main sequence star. If a protostar has less than 0.08 solar masses, it will never reach the temperature required for nuclear fusion (i.e., never join the main sequence). It will end life as a brown dwarf, which is a structure similar to Jupiter (but heavier). Protostars heavier than 0.5 solar masses will leave the Hayashi track after as little as a few thousand years to join the Henyey track.

On the Henyey track, the cores of these heavier protostars are hot enough for their radiation to punch through the outer layers, prompting a return to radiative heat transfer, decreased opacity, and a steady increase in luminosity. Consequently, the surface temperature of the protostar drastically increases as heat is effectively transported away from the core, prolonging its inability to ignite fusion in the core. However, the reduced heat in the core allows it to contract further under gravity, increasing its density and thus, it's temperature, once more. Eventually the heat reaches the level required to commence nuclear fusion.

Like the Hayashi track, protostars remain on the Henyey track for a few thousand to 100 million years, although heavier protostars remain on the Henyey track longer due to radiative heat transfer.

How a Main Sequence Star Evolves Over Time

Once hydrogen fusion begins, all stars enter the main sequence at a position dependent on their mass. Larger stars enter at the top left of the Hertzsprung Russell diagram (see above), while smaller red dwarfs enter at bottom right.

During their time on the main sequence, stars larger than the Sun will become hot enough to fuse helium. The inside of the star will form rings like a tree (see below), with hydrogen being the outer ring, then helium, then increasingly heavier elements toward the core (up to iron) depending on its size and temperature. These large stars remain on the main sequence for only a few million years, while the smallest stars remain for perhaps trillions. The Sun will remain for 10 billion years (its current age is 4.5 billion).

Fusion shells within a massive star. At the center is iron (Fe). Shells are not to scale.

Fusion shells within a massive star. At the center is iron (Fe). Shells are not to scale.

How a Main Sequence Star Begins to Die

When stars between 0.5 and 10 solar masses (e.g., the Sun) begin to run out of fuel, they leave the main sequence, becoming red giants. Stars larger than 10 solar masses typically destroy themselves in supernova explosions before the red giant phase can fully proceed. As previously described, red giant stars become particularly luminous due to their increased size and heat generation following the gravitational contraction of their cores. However, as their surface area is now much larger, their surface temperature decreases substantially. They move towards the top right of the Hertzsprung Russell diagram (see below).

A Hertzsprung Russell diagram that tracks the evolution of the Sun after it leaves the main sequence. (Image adapted from source.)

A Hertzsprung Russell diagram that tracks the evolution of the Sun after it leaves the main sequence. (Image adapted from source.)

As the core continues to contract toward a white dwarf state, the temperature may become high enough for helium fusion to take place in the surrounding layers. This produces a "helium flash" from the sudden release of this energy, heating the core and causing it to expand. The star briefly reverses its red giant phase as a result. However, the helium surrounding the core is quickly burnt, causing the star to resume the red giant phase.

How a Star Dies

Once all possible fuel is burnt, the core contracts until it becomes degenerate or a singularity, becoming super-hot in the process. Cores of less than 1.4 solar masses become white dwarfs (see below), which slowly cool to become black dwarfs. When the Sun becomes a white dwarf, it will be compressed to the size of the Earth and have about 60% of its mass.

Can you see Sirius A's tiny white dwarf companion, Sirius B? (lower left)

Can you see Sirius A's tiny white dwarf companion, Sirius B? (lower left)

Cores heavier than 1.4 solar masses (Chandrasekhar limit) will be compressed into 10-20 km wide neutron stars, and cores greater than approximately 2.3 solar masses (Tolman-Oppenheimer-Volkoff limit) will become black holes. It is possible for these objects to subsequently absorb enough matter to exceed these limits, prompting a transition to either a neutron star or a black hole.

In all cases the outer layers are completely expelled, forming planetary nebulae in the case of white dwarfs, and supernovae for neutron stars and black holes.

© 2013 Thomas Swan

Comments

Grace Marguerite Williams from the Greatest City In The World-New York City, New York on March 12, 2015:

Simply loved this hub. This hub is very well researched and contains great information. I have always been interested in stars and its evolution. Congratulations for being selected HOTD!

Ronald E Franklin from Mechanicsburg, PA on March 12, 2015:

This was both readable and interesting, Thomas. And I'm really glad to see that a hub that provides serious information on a scientific subject can be selected as HOTD. Congratulations!

Nico from Ottawa, ON on March 12, 2015:

very interesting. Love astronomy.

Anna Richmond from Fort Wayne, Indiana on March 12, 2015:

Very interesting! Nice job!

Thomas Swan (author) from New Zealand on November 09, 2014:

Thanks, glad you liked it. I've done likewise with a link in my "more information" section.

Doug West from Missouri on November 09, 2014:

I liked you Hub so I linked it to my Hub "Solar Flares and Their Impact on the Earth". Keep up the good work.

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