Articles will be posted and updated throughout the year, so feel free to return and check for updates.

Space and Time (coming soon)
-Interesting Objects
    Supernovae
    -Pulsars
    -Black Holes
-HR Diagram/Stellar Classification/Stellar Evolution
-Methods
    -Space Missions
        -mention furthest distance
    -Imaging
    Spectroscopy
    -Radio Astronomy
    -Photometry
    -Binary Stars
-Theories
    -Newton's Theory of Gravity
    -Special/General Relativity
    -Quantum Mechanics
    -String Theory
-links/books for more information

 

The term supernovae refers to a cataclysmic explosion that results in the death of a star or a stellar remnant.

There are actually several different types of supernovae. These types differ depending on what kind of star blows up. There are two main categories: Type I and Type II. Astronomers distinguish between these by their spectra (see spectroscopy); Type IIs spectra have hydrogen lines, while Type Is do not.

The divisions keep going from there: Type Ia, Type Ib, Type Ic . . . But that’s all pretty technical. The two most commonly discussed in astronomy are Type Ia and Type II.

To explain supernovae, though, we have to start with stars themselves. Stars are almost entirely hydrogen. Their second most common element is helium (the Sun is about ~ 70% hydrogen, ~28% helium). A star shines because it "burns" hydrogen — that is, the star fuses hydrogen into helium (via a chain process) in its core. Fusion, as we know from the Manhattan project and nuclear plants, produces a ton of energy, and it is this energy we feel and see from our star, the Sun. But a star only has so much hydrogen to fuse, and eventually the whole core will become helium. Fusion will stop.

Here’s where it gets interesting. There’s something called radiation pressure, which basically means that light/energy pushes on things. The energy radiated from a star’s core pushes outward on the star’s layers. This outward push balances the inward pull of the star’s own gravity upon those same layers (gravity wants to pull things to the center of mass — that’s why you don’t go flying off the Earth’s surface). Once hydrogen fusion stops, this outward pressure disappears. The star starts falling in on itself. The layers’ material falls into a smaller and smaller space, piling up, and eventually things become so dense and the temperature skyrockets so much that the helium core actually starts fusing. 

There can also be fusion in the shells of material directly outside the core, but that's still really far inside.

How many different rounds of fusion a star can undergo depends on its mass.  For small stars like the Sun, after helium fusion produces carbon the star will not be able to produce enough pressure to keep fusing. These stars end as balls of very dense carbon (a non-shiny diamond, if you will) spending the rest of existence cooling off until they’re just a cold orb in space. These are called white dwarfs.

White dwarfs are key to Type Ia supernovae. There’s a certain mass limit called the Chandrasekhar limit beyond which white dwarfs are no longer stable (for the curious, it’s 1.4 solar masses). Astronomers think Type Ia occur when a white dwarf with a larger, normal companion star siphons off material from the companion’s outer layers. This material piles onto the white dwarf, raising its mass past the Chandrasekhar limit. The white dwarf will try to collapse more because of this extra mass, but it’s already packed at electron degeneracy pressure, with all the atoms’ electrons packed together as tightly as possible. Unable to collapse further, the white dwarf explodes as a Type Ia. 

On the other side, Type II occur when much larger stars explode. If a star is at least several times as massive as the Sun, it will have enough core pressure (and a high enough temperature) to fuse carbon. A cycle ensues, from carbon to nitrogen to oxygen . . . if a star is massive enough, it will create iron. Iron is the last element possible to create in stellar fusion: to fuse iron into a heavier element requires putting more energy into the process than the fusion creates — unlike all the elements before iron, which produce more energy when fused than was needed to make them fuse.

Radiation pressure once again disappears. The layers fall inward, but the iron core, like the carbon core in the white dwarf, won’t fuse. As the outer layers are falling inward due to gravity, the core atoms, because of intense pressure, break down into neutrons, releasing a huge amount of energy. This energy blows out through the star’s outer layers; the infalling layers also rebound off the dense core. This combination creates a gigantic explosion called a Type II supernovae, throwing material into space at incredible speeds and temperatures. These explosions may leave behind a neutron star, the core of neutrons; a black hole, if the star is particularly massive; or no core at all, if the star manages to blow itself completely to bits. 

 Several supernovae events have occurred in Western civilization records, including Tycho’s Supernovae, which Tycho Brahe saw in the sixteenth century, and SN1987A, which astronomers spotted the night of 23 February 1987.

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Spectroscopy is the measurement of spectral lines from different atoms and molecules. In an atom, electrons can orbit the nucleus at different energy levels, which are unique not only for each element, but for each “species” of element — that is, for those with extra or fewer protons, neutrons, or electrons.

To move between these levels, an electron must swallow (“absorb”) or spit out (“emit”) the exact amount of energy that separates the two levels. These amounts are called quanta, hence the “quantization of energy,” and are absorbed or emitted as photons. To drop a level or levels, an electron must emit a photon (i.e. lose energy); to jump a level — or escape the atom completely — the electron must absorb a photon (i.e. gain energy). These photons have a unique wavelength, corresponding to the energy difference. We observe these wavelengths as spectral lines.

You’re likely familiar with the rainbow produced when sunlight shines through a prism. This rainbow is what astronomers (and physicists) call a continuous spectrum, because there are no clear separations in the spread. But light from a single element shows only a few lines instead of this continuum.

 

So why does all this stuff matter to astronomers? Since each line corresponds to a particular element, astronomers can determine which elements exist in space. Spectral lines reveal the compositions of stars, as well as of the interstellar medium — the gas and dust that lies between the stars. Spectroscopic observations also give information about the temperature of the atoms emitting or absorbing them, because certain temperatures provide the necessary energy to explain which lines astronomers see.

The kind of line seen also tells a lot about the region observed. Emission lines come from hot regions, like the gas around hot newborn stars. Absorption lines, on the other hand, come from cooler gas that lies between us and a hotter object. This gas could be the atmosphere of a star, which is cooler than the star’s inner layers.

Spectral lines are particularly important for stars. Astronomers separate stars into several types that depend on their light output and temperature, among other things, but since we’re no where near close enough to stick a thermometer in another star (besides our Sun), astronomers depend on spectral lines. Depending on which ionization levels — i.e. how many electrons have escaped from an atom — exist, astronomers can determine the star’s temperature. From that, they can work toward determining what kind of star they’re dealing with.

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