Category Archives: ISM Background


Stars will continue to forge elements into heavier elements in their cores via nuclear fusion until it is no longer energetically favorable for them to do so. The heaviest element that massive stars can create via nuclear fusion is iron. Once the core of a massive star has been entirely converted to iron, nuclear fusion cannot proceed further and there is no source of energy great enough to keep the star from collapsing under its own gravity. As a result, the outer layers of the star fall inward on the core and rebound, flying outward in a massive explosion called a supernova. Elements heavier than iron can be formed during this violent release of energy. The supernova generates huge shockwaves that propagate through the surrounding ISM, sweeping up material in front of them and heating and ionizing gas. Material from the outer layers of the star is also ejected into the ISM, enriching it in the heavier elements produced by nuclear fusion over the star’s lifetime. The next generation of stars to form in this region will therefore start out slightly richer in these heavy elements than their predecessors — so by measuring the amount of heavy elements present in stars, astronomers can infer the existence of prior generations of massive stars. All of the elements on Earth besides hydrogen (and some helium) were formed and released in this manner — so it is literally true that, as Carl Sagan said, “we are made of starstuff.”


Nuclear Fusion

Stars are powered by nuclear fusion, which is the process of smashing atomic nuclei together to create heavier elements. Stars like our Sun produce most of their energy by converting hydrogen to helium in their cores. This process releases energy in the form of light, which causes the stars to shine brightly. Stars like our Sun have enough hydrogen fuel to burn for about 10 billion years.

Stars more massive than our Sun also undergo nuclear fusion, but since they have more mass their cores are hotter and denser, and nuclear fusion proceeds more rapidly. After they convert all of their hydrogen to helium they begin to fuse the helium into carbon and other heavier elements, and so on. Each successive phase of fusion requires a higher temperature to proceed, so only the most massive stars will produce all of the heavier elements. Fusing heavier elements produces less energy per reaction each time, and producing elements heavier than iron actually requires more energy than it produces. Therefore, no matter how massive the star is nuclear fusion stops once iron is produced, and the star will then collapse in a supernova explosion.

Understanding Spectral Lines

The modern picture of the atom is a compact nucleus, containing protons and neutrons, surrounded by a cloud of electrons. These electrons can only occupy specific energy states. If a photon with exactly the right amount of energy hits an atom, it can push an electron from a lower energy state to a higher energy state, or even knock the atom free of the nucleus entirely (a process called ionization.) Electrons naturally like to be in the lowest possible energy state, so if an electron is bumped into a higher (“excited”) energy state it will eventually spontaneously fall back to a lower energy state on its own. When it does, it will emit a photon equal to the energy difference between its starting state and its final state. We can see evidence of both steps of this process: if we are looking through cool gas with stars behind it, the gas will absorb the starlight as electrons in the atoms in the gas are excited, so we see dark lines in the spectrum from the star corresponding to these transitions. If, instead, we are looking at a bright object we will see the emission released when the electrons fall back down to lower energy states, so we will see bright lines at the points in the spectrum corresponding to these transitions.

Since the energy levels of different elements are different, if astronomers see a specific spectral line coming from a certain region, they can infer what elements are present there. One particularly important line is the hydrogen-alpha line at 656.28 nm. This is caused by the transition of an electron in a hydrogen atom from the third lowest energy state to the second lowest. This transition frequently results when an ionized hydrogen atom recombines (i.e. recaptures an electron to become neutral), so H-alpha emission is characteristic of regions of ionized hydrogen. The H-alpha emission from the entire sky has been mapped in great detail, revealing a wealth of structure in our galaxy.

Blackbody Radiation

A blackbody is an object that absorbs all electromagnetic radiation that hits it and re-emits it at a range of wavelengths determined solely by its temperature. The shape of a blackbody’s spectrum is always the same, but the peak of the spectrum shifts to shorter wavelengths (higher energies) as the blackbody’s temperature increases. (This is why a cold piece of metal at room temperature appears dark, but if heated in a fire it will begin to glow red, then orange, and then white as the peak of its emission shifts to shorter wavelengths.)

The theoretical blackbody spectrum for an object at various temperatures. Hotter objects emit more total radiation (i.e. they are brighter) and the peak of their spectrum shifts to shorter wavelengths as temperature increases
The theoretical blackbody spectrum for an object at various temperatures. Hotter objects emit more total radiation (i.e. they are brighter) and the peak of their spectrum shifts to shorter wavelengths as temperature increases.

It turns out that stars are well-approximated as blackbodies. Our Sun, for example, has a temperature of 5800 K, which puts the peak of its spectrum at about 500 nm — which is the wavelength of green light. It makes sense that our eyes evolved to see a narrow range of wavelengths corresponding to the peak of the solar spectrum!

Since the peak wavelength of a blackbody spectrum is inversely proportional to its temperature, all astronomers have to do to obtain the temperature of a glowing object is to  measure its spectrum and determine the peak wavelength. This is incredibly useful for studying stars. Another tool that is important for stars and the ISM alike is the study of spectral lines.

The Electromagnetic Spectrum

When sunlight passes through a prism, we see it split into a rainbow of colors. This effect can be explained by understanding light as an electromagnetic wave. (We therefore sometimes refer to light as electromagnetic radiation.) As the light passes through the prism, shorter wavelengths are bent more than longer ones, which produces the familiar rainbow. Longer wavelengths (red) have lower energy than shorter wavelengths (blue). We call a graph of how much light an object emits at every wavelength a spectrum.


Our eyes are fine-tuned to see only the tiny range of wavelengths that are most strongly emitted by our Sun. (The amount of light emitted at each wavelength is largely determined by the physics of blackbody radiation.) However, there are many kinds of light that we can’t see, spanning a wide range of wavelengths. A diagram of the full electromagmetic spectrum is shown below. The lowest energy light has wavelengths as big as mountains, while the highest energy light has wavelengths as small as the width of atoms. Astronomical processes produce light at all of these different wavelengths. By looking at the Universe with telescopes that can see different portions of the electromagnetic spectrum, we can see different processes occurring in the same region of the sky and gain new information about what is happening there. For example, hot young stars produce a lot of light in the ultraviolet and X-ray portions of the electromagnetic spectrum while older cooler stars do not. Additionally, while dust in the ISM can only be seen in visible light when light from stars reflects off it or backlights it, it glows quite brightly at infrared and radio wavelengths.

Diagram of the electromagnetic spectrum, showing the wavelength (distance between consecutive wave peaks) and frequency (number of wave cycles per second) of different kinds of electromagnetic radiation, from radio waves to gamma rays. Note how small the portion of the electromagnetic spectrum that we can see is.

What is dust?

The dust in the interstellar medium is not like the dust you might find in your house. Firstly, the size of the dust grains is much smaller — they are typically less than 0.1 \mu m in diameter (more than a hundred times smaller than the width of a human hair)! Secondly, it is much less dense. With only one particle per cubic centimeter, collisions between dust particles are very rare. Understanding exactly how dust grains eventually come together to condense into planets and stars is an ongoing area of active research.

Dust is found throughout molecular clouds and can be seen in beautiful images of star forming regions like the Orion Nebula. Typically, there is about one hundred times more gas in these molecular clouds than there is dust. However, the dust is crucial the formation of these clouds. In order to form a hydrogen molecule, two individual hydrogen atoms must pass sufficiently close together to form a bond. In the very low-density ISM, such close interactions are rare. However, dust grains, while tiny, are still much, much larger than individual atoms — so interactions in which a hydrogen atom condenses onto a dust grain are far more common. When two hydrogen atoms condense onto the same dust grain, they can move around the surface until they find each other, and then are able to bond together to form a molecule, which detaches from the dust grain and goes on its way. Without the presence of the dust grains to act as a catalyst, the conversion of hydrogen from an atomic state to a molecular state would occur much too slowly to match the observations we see of the ISM.

Many dust grains have similar compositions to comets and asteroids. They may be composed of silicates, carbon grains, or water ice. Carbon monoxide is an important tracer of dust density, although in the densest, coldest dust clouds it may freeze onto dust grains and become undetectable.

A Note about Density

Most of the ISM is extremely diffuse. The average density is just one particle per cubic centimeter! (In comparison, a typical cubic centimeter of air at sea level contains about 3\times10^{19} — that’s 30,000,000,000,000,000,000 — particles.) Even the densest regions of the ISM contain only about 1,000,000 particles per cubic centimeter — which is almost as close as modern lab equipment can get to a perfect vacuum. One interesting effect of these incredibly low densities is that particles don’t encounter each other very often, which gives individual atoms enough time to undergo improbable energy transitions between collisions, producing “forbidden line” emission. These “forbidden” spectral lines are important to astronomers because they help reveal the composition of the ISM.

What is the ISM?

The interstellar medium, or ISM, is the formal term for everything that fills the spaces between the stars in a galaxy. The ISM of the Milky Way and other nearby galaxies has been extensively studied by astronomers. It consists mostly of hydrogen, along with other gases and dust grains. Some of this material will eventually coalesce to form stars. When these stars reach the end of their lives, the material will return to the ISM, enriched by nuclear fusion processes. Scientists can therefore study the ISM to learn how stars form and how previous generations of stars in the galaxy lived and died.

The Orion Nebula, part of the Orion Giant Molecular Cloud and a site of active star formation.
The Orion Nebula, part of the Orion Giant Molecular Cloud and a site of active star formation. Dust and ionized, atomic, and molecular hydrogen all exist in the nebula.

As we will see on the rest of the tour, the ISM consists of mostly hydrogen gas. We can divide the ISM into “phases” based on its temperature, density, and ionization state:

  • Coronal Gas (Hot Ionized Medium, HIM): Very hot (1,000,000 K), very diffuse ionized gas, heated by supernovae and other shocks.
  • HII Gas: Ionized hydrogen at temperatures of about 10,000 K. Either “diffuse HII” if not surrounded by a molecular cloud, or “HII region” if bounded by neutral gas.
  • Warm and Cool HI (Warm and Cool Neutral Medium, WNM and CNM): neutral atomic hydrogen at “warm” temperatures (5,000-10,000 K) and “cool” temperatures (50-100 K). The existence of the WNM is debated.
  • \textbf{H}_2: cold (17-20K) molecular hydrogen, protected from ionizing radiation by a surrounding layer of neutral atomic hydrogen.
  • Dust: very small solid particles, typically less than 0.1 \mum. The composition of interstellar dust is discussed here.

All of these phases are observed on the tour, with the exception of the coronal gas, which is too diffuse to easily observe in our own galaxy. Often, multiple phases can exist at the same time; with no sharp boundary between, say, regions of molecular hydrogen gas and regions of dust.

The ISM also includes dark matter, photons (light) from stars and other sources, cosmic rays (particles that have been accelerated to extremely high speeds), and gravitational and magnetic fields. All of these things can interact with the gas and dust we see, contributing to the rich complexity of the ISM as seen in images like the Hubble Space Telescope photo above.