Starstuff: The Cosmic Origins of Elements and Supernova Chemistry

Spectroscopy allows scientists to determine what stars are made of by analyzing their light across visible and other wavelengths without physically visiting them. Stars contain small amounts of almost all chemical elements, and together they produce thousands of absorption lines in their spectra. By matching these lines to known emission spectra of elements, astronomers can identify the chemical composition of stars and other celestial bodies.

In 1802, William Herschel discovered that sunlight produces a spectrum with dark lines. In 1814, Joseph von Fraunhofer identified hundreds of these lines, now called Fraunhofer lines. Each element produces a unique pattern of spectral lines, like a fingerprint. By matching observed spectral lines to known patterns, astronomers can identify the elements present in stars. This technique has revealed that the Sun and other stars contain elements like hydrogen, helium, iron, sodium, and many others.

The spectroscope works by dispersing light into a spectrum of colors from red to violet. The Sun's spectrum is riddled with fine dark lines because light is a wave and the lines mark particular wavelengths where some of the energy in the Sun's light has gone missing. In the 1850s, two German physicists Robert Bunsen and Gustav Kirchhoff discovered that when light passes through a gas, the atoms in the gas absorb some of the light but only at certain specific wavelengths, which depend on the properties of that particular type of atom. This explains the lines in the Sun's spectrum, called absorption lines. Kirchhoff and Bunsen realized that the lines act as a chemical fingerprint because they reveal which gases are present in the Sun's atmosphere. However, this is only half the story. When atoms are energized in the right way, they can emit light too, creating a spectrum of bright lines known as emission lines, which are the counterparts to the dark absorption lines. For astronomers, all of these spectral lines—both bright and dark—revealed a link between Earth and Sky. In the laboratory, they could measure the spectra of gases and flames using known materials, then compare these results to the spectra of the Sun and stars. Starting in 1862, the English astronomer William Huggins was among the first to attach a spectroscope to his telescope and explore the heavens in a new way. He was later joined by his wife and scientific collaborator Margaret. From their home-built observatory near London, Huggins recorded the spectra of many stars and other celestial objects. Among his targets was the great nebula in Orion, a misty patch of light located near the stars of Orion's belt. Huggins found that the nebulous spectrum consisted entirely of bright emission lines, especially lines corresponding to hydrogen. Here was evidence that the nebula was really a vast cloud of glowing gas, perhaps a cloud that might someday condense to form new stars. In 1864, Huggins saw something a bit different when he turned his spectroscope on the Cat's Eye Nebula, an expanding shell of gas blown off of a dying star. The Cat's Eye produced a bright green emission line that didn't seem to correspond to any known element. Huggins dubbed the mystery substance nebulium. It would take over 60 years before astronomers realized that the green line is produced by atoms of oxygen with two electrons stripped away. This energized state reveals the high temperature and low pressure environment within the nebula, which is expanding at nearly 60,000 kilometers per hour. In 1868, the spectroscope really did lead to the discovery of something entirely new. In August of that year, French astronomer Jules Janssen used a spectroscope during a total eclipse of the Sun, hoping to isolate the light from a solar prominence—a glowing ribbon of ionized gas suspended in the Sun's powerful magnetic field. Normally prominences are impossible to see because the surface of the Sun is so bright, but during a total eclipse, any large prominence is on the edge of the Sun and immediately visible to the eye. When Janssen turned his spectroscope on one of them, he saw some emission lines he didn't recognize. Like Janssen, Norman Lockyer in England had also seen the emission lines. Both recognized that the lines were produced by a previously unknown element. Lockyer named it helium after Helios, the Greek god of the Sun. It would take 27 more years before the existence of helium was finally confirmed in the laboratory, but its discovery on the Sun made Janssen famous. By the end of the 19th century, the spectroscope had proved its worth in decoding the chemistry of the Sun and the stars. But its greatest triumph was still to come. By 1925, it seemed that the quest to understand what stars are made of had come full circle. 24 centuries earlier, Anaxagoras had written that stars are fiery stones. Now the spectroscope was suggesting something similar. By pointing it at the Sun and distant stars, astronomers could see spectral lines that corresponded to sodium, calcium, magnesium, iron, and many other elements that are commonly found in minerals on Earth. This persuaded the American astronomer Henry Norris Russell to suggest that if Earth's crust were heated up to the Sun's temperature, it would share the Sun's spectrum. Russell was a leading authority on stars and his view defined the scientific mainstream. But Russell's Universe was about to change.

Astronomers use stellar spectroscopy to analyze star composition by splitting light into its component wavelengths. Each spectral line corresponds to a different element or molecule, acting like cosmic IDs. The behemoth star shows complex mixtures including silicon monoxide, carbon monoxide, water vapor, and more complicated molecules forming in its dust shell. This reveals the chemical processes occurring in extreme stellar environments.

This segment traces the fundamental chemical evolution of the universe. It begins with the astronomical definition of metallicity, where hydrogen and helium are non-metals while all heavier elements are classified as metals. The video explains how the universe started with mostly hydrogen and helium, then gradually became enriched with heavier elements through stellar processes and supernovae. Stars are classified into three populations based on metallicity: Population I (youngest, richest in metals), Population II (older, fewer metals), and Population III (the first stars, containing no metals). These first stars, formed before any supernovae occurred, are expected to be super massive (hundreds of solar masses), extremely bright blue stars that would exist for only 1-2 million years before exploding. This cosmic chemical progression forms the foundation for understanding stellar and galactic evolution throughout the universe's history.

Stars are classified into three generations based on their metal content: Population III stars (theoretical first generation) formed from primordial hydrogen and helium after the Big Bang; Population II stars are older and metal-poor, formed from gas polluted by earlier supernovae; Population I stars (including our Sun) are younger and metal-rich, formed from gas enriched by multiple generations of stellar explosions. Our Sun contains mostly primordial hydrogen and helium with trace metals from various sources including red giant stellar winds, neutron star collisions, and supernovae. The metallicity of a star depends on when and where it formed in the universe's history.

Stars are classified into populations based on their age and metallicity. Population III stars were the first stars, formed from primordial hydrogen and helium, existing briefly before exploding as supernovae. Population II stars are second-generation stars formed from gas enriched by Population III supernovae, characterized by extremely low metallicity. Metallicity is measured by the iron-to-hydrogen ratio; our Sun has 1 iron atom per 20,000 hydrogen atoms, while metal-poor stars can have ratios 1,000 to 10,000 times lower. Finding Population II stars outside the Milky Way allows astronomers to compare galactic chemical evolution.

Metallicity refers to the percentage of a star composed of elements heavier than helium. Our Sun is classified as Population 1 with metallicity about 2% by mass (about 0.1% by atom count). Population 2 stars have lower metallicity, down to about one-thousandth of our Sun's, and are older stars, mostly less massive than our Sun. Population 3 stars are hypothetical first-generation stars with essentially no metallicity. The Sun formed about 9 billion years after the Big Bang with slightly under 2% heavier matter by mass.

Stars are classified into populations reflecting their ages and chemical compositions. Population I stars are young, metal-rich objects like our Sun, found predominantly in the galactic disk. Population II stars are older, metal-poor stars inhabiting galactic halos and globular clusters. Population III stars, hypothetical first-generation stars, would have zero metallicity, formed from primordial gas shortly after the Big Bang. HD 140283, estimated at 14 billion years old, approaches the age of the universe itself. This population structure reveals galactic chemical evolution, with each generation enriching the interstellar medium with heavier elements produced through stellar nucleosynthesis. The chemical composition of stars thus serves as a cosmic record of stellar evolution and galactic history.

Galactic chemical evolution models how the abundance of elements changes over cosmic time. Stars produce elements through nucleosynthesis, then return them to the interstellar medium through stellar winds or supernovae. These enriched materials form new stars and planets, creating successive generations with progressively heavier elements. Nuclear astrophysics experiments provide the reaction rates needed for these models.

Galactic chemical evolution models track how elemental abundances change over cosmic time through stellar birth, evolution, death, and recycling. Simple one-zone models require initial conditions from Big Bang nucleosynthesis, star formation rates, initial mass functions, stellar yields, binary fractions, and gas dynamics. More sophisticated models incorporate spatial variations and time-dependent processes. Comparing model predictions with observed abundances reveals systematic discrepancies—for example, underproduction of cobalt, potassium, and some s-process elements. These mismatches guide improvements in stellar evolution theory and highlight the need for better constraints on stellar yields and binary interactions.

Galactic chemical evolution models track how elemental abundances change over cosmic time. Initially, galaxies contain only hydrogen and helium from the Big Bang. As stars form and die, they produce heavier elements that enrich the interstellar medium. New generations of stars form from this enriched material, creating a progressive increase in metallicity over time.

Galactic chemical evolution traces how elements formed and distributed in the interstellar medium. Light elements formed during the Big Bang, while heavier elements were synthesized inside stars. The first stellar generations were metal-free, enriching subsequent generations through stellar death. Chemical evolution models use three fundamental terms: gas converting to stars, stellar yields of ejected elements, and gas inflow. The Initial Mass Function (IMF) describes stellar mass distribution at birth, traditionally ranging from 0.1 to 100 solar masses with a Salpeter slope of 1.35. Stellar yields depend on mass and metallicity, with AGB stars (M<8 Msun) producing carbon, nitrogen, and sodium, while massive stars (M>8 Msun) contribute alpha elements and iron through supernovae and winds. The final fate of stars depends critically on mass and metallicity: electron capture supernovae for 8-9 Msun stars, core collapse supernovae for M>11 Msun, and pair instability supernovae for very massive stars at low metallicities. These different death channels produce distinct elemental yields, with pair instability supernovae producing disproportionately large iron yields. Chemical evolution models require calibration against multiple observational constraints including metallicity distribution functions, star formation histories, supernova rates, gas fractions, and solar metallicity.

Chemical evolution models simulate the formation and enrichment of chemical elements in the Milky Way by tracking star formation, stellar nucleosynthesis, and interstellar medium recycling over cosmic time. These models predict elemental abundances in different galactic components (halo, thick disk, thin disk, bulge) and enable comparison with observational data from spectroscopic surveys. The models incorporate processes such as supernova enrichment, stellar migration, and gas dynamics to explain how the chemical composition of our galaxy has evolved from its primordial state to the present day.