The death of a star

As T.S. Eliot probably meant to say, “This is the way a solar-mass star ends / Not with a bang but with a whimper.” When a star like our sun runs out of fuel, it eventually puffs up and releases its outer layers, leaving behind a dense inner core called a white dwarf. In isolation, this white dwarf isn’t particularly interesting. It’s no longer fusing any nuclei together to produce energy, so it just slowly cools and dims over time.

But every once in a while, a white dwarf explodes, producing a Type Ia supernova (Figure 1). (Note: the classification “Type Ia” is, like most nomenclature in astronomy, a historical artifact that no one’s bothered to change even though it’s now extremely confusing.) Type Ia supernovae are incredibly useful to astrophysicists because their light curves are standardizable, meaning that they can be used to measure distances using the “standard candle” method—that is, because we know exactly how intrinsically bright a supernova is, we can figure out how far away it is based on how bright it appears to be (since supernovae that are farther away will look dimmer). With this technique, Type Ia supernovae were used to measure the accelerating expansion of the Universe, a Nobel Prize-winning discovery!

Figure 1. This is one example of what’s left behind after a Type Ia supernova. This beautiful image is an X-ray and infrared composite of supernova remnant G299.2-2.9. Credit: Chandra/Spitzer

How do you solve a problem like an exploding white dwarf?

However, despite the importance of Type Ia supernovae, we don’t actually understand how they work. What exactly makes white dwarfs explode?

The classic “textbook” theory about Type Ia supernovae is fairly straightforward. As mentioned earlier, an isolated white dwarf is pretty boring… but if it has a companion, things can get a lot more interesting. The gravitational pull of the white dwarf can disrupt the companion star (traditionally a massive red giant star), pulling some material from the companion onto the white dwarf. In this way, the white dwarf can accrete matter onto itself (Figure 2) and grow in mass.

Figure 2. An artist rendition of a white dwarf (right) accreting material from its companion until it explodes. Credit: NASA/CXC/M Weiss

But the white dwarf can’t grow forever. At some point, it reaches a limit of approximately 1.4 solar masses called the Chandrasekhar limit; above this limit, the quantum degeneracy pressure supporting the white dwarf is overwhelmed, and the star quickly explodes, producing a supernova.

Simple enough, right? Except this model doesn’t seem to explain all the observations about white dwarfs and Type Ia supernovae. For instance, we don’t see many white dwarfs close to the Chandrasekhar limit, and simulations don’t produce the same amount of heavy elements that we measure in supernova remnants.

To solve these problems, astrophysicists have come up with several other potential mechanisms that could also cause Type Ia supernovae. Most of these involve some variation on the textbook picture of a “fast explosion of a single white dwarf at the Chandrasekhar limit.” Here are a few of the more common theories:

  • Slow, instead of fast: If the white dwarf doesn’t explode immediately, it has some time to expand, which allows it to more easily produce the kinds of elements that we measure in observations. This is called the “deflagration-to-detonation transition.”
  • Sub-Chandrasekhar instead of Chandrasekhar: A white dwarf below the Chandrasekhar limit can explode if it can accrete a lot of helium from its companion star. The helium builds up on the white dwarf’s surface, eventually causing a runaway nuclear reaction that sets off an explosion. This is called the “double detonation” scenario.
  • Double instead of single: Instead of one white dwarf, two white dwarfs (usually both below the Chandrasekhar limit) can merge or collide, causing an explosion. This is called the “double-degenerate” channel.

It’s likely that Type Ia supernovae are not caused by any single one of these scenarios, but by a combination of some or all of them. But how can we tell which mechanisms are the most common?

Manganese? I’m a fan-ganese

This is where galactic archaeology comes in! Supernovae produce lots of heavy elements, and as Evan’s post describes, we can use ratios of different heavy elements to deduce information about past explosions in galaxies. The example in Evan’s post looks at magnesium (Mg), which tells us about the fraction of heavy elements produced by Type II supernovae (explosions caused by the collapse of high-mass stars).

To learn more about Type Ia supernova, though, we can look instead at other elements like manganese (Mn). Manganese, like iron, is mostly produced by Type Ia supernovae (again, see Evan’s post)—but it’s incredibly neutron-rich, with 5 more neutrons than protons. In order to produce large amounts of such a neutron-rich element, the original white dwarf must be very dense. This is true for some of the models described above (like the “textbook” model and the “deflagration-to-detonation transition” model) but not for others (like the “double-detonation” and “double-degenerate” models).

So if we can measure the amount of manganese produced by Type Ia supernovae, we can distinguish among some of these models!

This is what I’m working on now. I’m taking spectra using the DEIMOS instrument on the Keck Telescope to measure manganese in stars in nearby dwarf galaxies. Figure 3 shows an example of one of these spectra, where some of the absorption lines caused by manganese are shaded.

Figure 3. Part of an observed spectrum from a star in the Sculptor dwarf galaxy (black points). The shaded regions show locations of absorption lines produced by manganese; different manganese abundances produce different synthetic spectra (colored lines). The best-fitting line (orange) is the spectrum with a manganese abundance [Mn/H]=-1.58.

These are moderate-resolution spectra (meaning they’re a little more “blurry” than high-resolution spectra, but they take less time to measure), so it’s difficult to actually measure the amount of manganese from any one line. Instead, I can create synthetic spectra using models of stellar atmospheres. Different manganese abundances produce different synthetic spectra, which are shown as colored lines in Figure 3. By figuring out which spectra best fits the data, we can figure out precisely how much manganese a star has. In the figure, it’s pretty clear that one of these models is a better match for the data than the others: the one with [Mn/H] = -1.58 (this notation just means that this star has 10-1.58=0.026 times as much manganese as the Sun).

I’m still working on doing this kind of analysis for several stars in nearby dwarf galaxies. Check back soon to see the next step: using manganese measurements from lots of different stars, we can figure out how much manganese is produced by Type Ia supernovae. And eventually, this galactic archaeology will tell us something about the physics of Type Ia supernovae!