The most powerful way a star can die is through a supernova explosion. Through nuclear physics and the pull of gravity, a star can be ripped apart, shining brighter than an entire galaxy for a brief moment of cosmic time. There are several kinds of supernovae, but one particular kind known as Type Ia is particularly interesting to astronomers. And we still aren’t entirely sure how they occur. 

Type Ia supernovae are useful to astronomers because they always explode with about the same intensity. By measuring the apparent brightness of distant Type Ia supernovae, we can calculate the size of our universe, and how rapidly it’s expanding. But what makes a Type Ia supernova so consistent? Other supernovae can vary in brightness, so why not these?

The most popular theory is that they are caused by particular kind of binary star, where a white dwarf orbits a red giant. White dwarf stars typically have a mass about equal to our Sun. That’s because if they get much larger they reach what is known as the Chandrasekhar limit, and collapse under their own weight. If a white dwarf happens to orbit closely to a red giant, outer layer of the red giant can be captured by the gravitational pull of the white dwarf. This increases the mass of the white dwarf. And if its mass reaches the limit, then boom!

Evidence of an early helium flash before a supernova is triggered. Credit: Institute of Astronomy, the University of Tokyo

At least that’s the idea. The challenge is to work out the details. Just how does the gravitational collapse of a white dwarf trigger the type of nuclear reactions necessary to cause the star to explode? But new research shines more light on that question.

Using data from the Subaru telescope and the Gemini-North telescope, a team of astronomers captured the earliest moments of a Type Ia supernova. They found that a small initial flash occurred before the much larger supernova explosion. The spectrum of the initial explosion seems to indicate that it was caused by the rapid fusion of helium on the star’s surface. This effect has also been seen in computer simulations.

In this model, a build up of helium on the surface of the white dwarf causes it to start fusing rapidly. The explosive detonation of helium creates shock waves that propagate to the center of the star, where they trigger a chain reaction of fusion in its core. In this way, the explosion is not caused directly by reaching the Chandrasekhar limit, but by a chain reaction of rapid nuclear fusion. This would explain why there is some small variation in the brightness of Type Ia supernovae. It could also help us better refine how these supernovae are used to measure galactic distances.

We’ll need to observe more of these supernovae to learn whether this is the primary trigger mechanism for Type Ia supernovae, or if there are other ways these stars can be detonated. But for now it seems clear that with a bit of helium we can detonate a star.

Paper: Ji-an Jiang, et al. A hybrid type Ia supernova with an early flash triggered by helium-shell detonation. Nature 550, 80–83 (2017)

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Our home planet is exceptional for being a warm rocky planet with plenty of water on its surface. There have been several proposed origins for Earth’s water, such as that it originated within the primordial materials of our planet, or that it was brought by cometary or meteor impacts. We know from measuring isotopic ratios in Earth’s water that much of it was formed before even our solar system, and that it seems to match water found in certain asteroids. So the most popular model for Earth’s water is that it was brought to our planet by asteroid-like meteors. But did asteroids in our young planetary system really have enough water to produce our oceans?

It would seem the answer is yes, and new research tends to support that conclusion. In this work, the authors looked at the spectra of a white dwarf and found evidence of hydrogen at oxygen in its atmosphere. A white dwarf is a sun-like star that has reached the end of its life. After is has consumed much of its hydrogen, fusing it to heavier elements like helium and carbon) it collapses under its own weight until it is roughly the size of Earth (but still the mass of the Sun). Because a star will enter a red giant stage before collapsing to a white dwarf, most of the lighter elements like hydrogen would be cast off. Likewise, heavier elements like oxygen would tend to settle into the core of the star. So one would think we wouldn’t see much of either hydrogen or oxygen in the spectra of a white dwarf.

It turns out we do tend to see hydrogen, which could be an indication that some of that light element didn’t get thrown off during the red giant stage. But the presence of oxygen would seem to indicate this material was accreted by the star relatively recently on a cosmic scale. The authors found evidence of other elements such as silicon and iron, which are common in asteroids. One explanation for this is that the white dwarf accreted an asteroid, and its remnants are seen in the star’s spectra.

So the authors calculated the size of such an asteroid from the spectra, and found that it would have been about the size of Ceres. If that’s the case, then the strength of the oxygen and hydrogen spectra would imply that the asteroid was about 38% water when it was accreted. That’s a sizable amount. We know that asteroids in our solar system such as Ceres and Vesta contain water, but the amount is still under investigation. If this roughly 1/3 fraction is common, then there was plenty of water for a young Earth to gain by meteor impacts. The fact that this was seen around a white dwarf is also an indication that water-bearing asteroids may be common in planetary systems.

This isn’t conclusive proof that Earth’s water did in fact come from meteor impacts, but it adds to a growing pool of evidence supporting that model.

Paper: R. Raddi, et al. Likely detection of water-rich asteroid debris in a metal-polluted white dwarf. MNRAS 450 (2): 2083-2093 (2015)

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A recurrent nova is a stellar nova that occurs from the same star more than once.  You’re probably more familiar with a supernova, where a star is ripped apart in a cataclysmic explosion.   Novas occur when a white dwarf orbits with another star and captures some of the star’s outer material. This material forms an accretion disk around the white dwarf, which gradually falls to its surface.  When material accumulates on the surface of the white dwarf, it can trigger a nuclear explosion that causes it to brighten similar to a supernova, but not nearly as intense.  Since the explosion doesn’t destroy the star, it is possible for a nova to occur again after more material has accumulated.

The most famous recurrent nova is RS Ophiuchi, which becomes a nova about every 20 years.  Other recurrent nova occur at different rates.  Recurrent nova caused by more massive white dwarfs tend to have shorter repeat times than those with less mass.  This seems to be due to the fact that stars with more gravity can accumulate matter from the companion star more quickly.  Some recurrent nova have irregular repeat times, or have novae with highly varying brightnesses.  This seems to be due to disruptions in the accretion disk when the star explodes.

It is thought that recurrent novas could be a precursor to the white dwarf becoming a supernova.  If the material cast off by the nova is less than the material accumulated each time, then the white dwarf will gradually increase in mass.  Eventually this bring its mass to the Chandrasekhar limit, which is the upper limit for the mass of a white dwarf.  Beyond that point the star will collapse, which can trigger a supernova explosion.

Background: The star as a Nova.
Inset: Hubble image of the star when inactive.

Recently astronomers have discovered a fast recurrent nova in the Andromeda galaxy.  A paper on the nova was recently published in Astronomy and Astrophysics, and presents some of the initial results.  The star, known as M31N 2008-12a has a nova outburst about once a year.  This is extremely rapid, since the next highest frequency rate is about once a decade.  The star also brightens quickly and dims quickly, on the order of a few days.  When observed in x-rays, it was found the star emits x-rays during the nova period for about 10 days.  Since x-rays are generated by nuclear interactions, this indicates that the nuclear interactions on the surface only last about that long.

All of this points to M31N 2008-12a being a very massive white dwarf star.  Since the novae occur so rapidly, it is clear that material from its companion star continues to accrete at a regular basis.  It would seem that this star is on its way to becoming a supernova.  Just how soon that might be is unknown, but it’s worth keeping an eye on.  It could explode at any time.

Of course on a cosmic scale “any time” could be anywhere from tomorrow to thousands of years.

Paper: M. Henze, et al. A remarkable recurrent nova in M 31: The X-ray observations. Astronomy & Astrophysics, 563, L8 (2014)

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Last time I talked about how large stars can become a supernova through a collapse of their core.  But this only occurs in stars much larger than our Sun.  So how can a solar mass star become a supernova?  For that, it has to dance with another star.
If you recall from last time, when a Sun-sized star begins to run out of hydrogen to fuse in its core, it fuses helium for a while, causing it to swell into a red giant.  After this stage the star can’t fuse any higher elements, and what remains of the star compresses down to a white dwarf.  All the mass of the star is squeezed to about the size of the Earth.  In a white dwarf it isn’t the heat and pressure of fusion that balances against the weight of gravity, but the pressure of the electrons pushing against each other.  For a single star such as our Sun, that’s the end of the story.  The white dwarf simply cools over time.  But for a binary star system, something more interesting can occur.

Suppose there was a binary star system, consisting of two stars similar to the Sun, and reasonably close together.  These two stars will spend most of their lives as main sequence stars, but they will eventually run out of hydrogen and become red giants.

Of the two stars, one is likely to be a bit larger than the other.  This means throughout it’s main sequence period the larger star will burn just a little bit hotter, and use up its hydrogen a bit faster.  As a result, the larger star will enter its red giant stage while the smaller star is still a main sequence star.

A red giant is fairly large.  For a Sun-like star the red giant can be as large as the Earth’s orbit.  In a binary system this means the outer layers of the red giant can can be captured by the companion star.  It can even engulf the two stars in what is called a common envelope.  The result is that some of the mass of the larger star is transferred to the smaller star during this stage.  The larger star then collapses into a white dwarf.

Some time later, the smaller star enters its red giant stage, and a similar process occurs.  The smaller star swells into a large red giant, and some of the outer material is captured by the first star, which is now a white dwarf.  You can see an artist’s rendering of this process in the image above.

A white dwarf is formed because the electron pressure within the white dwarf is able to balance its gravitational weight.  But this electron pressure has a limit (about 40% greater than the mass of the Sun) known as the Chandrasekhar limit.  As the white dwarf continues to capture mass from the other star, its mass continues to rise, approaching that ultimate limit.

As the this limit is reached, the star begins to collapse into a neutron star.  This causes a cascade of rapid fusion within its core, which rips the white dwarf apart.  The white dwarf explodes as a supernova.  This stellar dance ends with a bang.

What makes this type of supernova particularly interesting is that it always happens with a white dwarf of about 1.4 solar masses, and it always occurs through the same general process, destroying the white dwarf completely.  This means a supernova created in this way always has about the same brightness.  They are called type Ia supernovae, and they are very useful to astronomers.

There are various ways to measure the distances of galaxies.  One way is to observe a particular kind of star that varies in brightness, known as a Cepheid variable star.  We can use these stars to measure the distance of closer galaxies, but only as far as about 90 million light years.  For more distant galaxies these stars are just too dim to observe accurately.  But we’ve also observed type Ia supernovae in galaxies for which we have measured their distance.  We can observe how bright the supernova appear, and knowing their distance we can determine how bright they actually are.  What we find is that type Ia supernova always have the same absolute brightness.

This means we can use them as a “standard candle”.  If we observe a type Ia supernova in a distant galaxy, we can observe how bright it appears.  Since we know how bright it actually is, we can calculate the distance to the galaxy, since the more distant a light source is, the dimmer it appears.  We can therefore use this type of supernova to measure the distance to its galaxy.

Since these supernovae are very bright, we can measure the distance of galaxies billions of light years away.  Without them, the size and structure of the universe would still be a great mystery.

Two stars, having a last dance, and helping us to discover the secrets of the cosmos.

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Last time I talked about how black holes can exceed a maximum luminosity known as the Eddington limit.  The Eddington limit is the point where a star is so bright that its light pressure balances out the gravitational pressure of the star.  Today I’ll talk about another limit known as the Chandrasekhar limit.  This one puts a limit on how massive a star can be when it dies, and it relies on a bit of quantum mechanics.

When a typical star is in its active lifetime (main sequence), the gravitational weight of its mass is balanced by the heat and pressure generated by fusion in its core.  This is a stable period for a star, and can last for billions of years (or in the case of red dwarfs, trillions).  But eventually a star runs out of hydrogen to fuse, and it can’t sustain the heat and pressure in its core.  With no more fusing core, the old star is basically just a hydrostatic mass. The key difference is that the density is so high it isn’t a regular ball of gas, but rather a gas of plasma. All the ions and free electrons act somewhat like a regular gas, but with one big exception: the electrons and ions obey the Pauli exclusion principle.

In broad terms the Pauli exclusion principle means that the electrons can’t occupy the same state. It is kind of like musical chairs where everyone jostles each other to find a seat, and two people can’t share the same seat. The exclusion principle also applies to the ions, but it turns out that the electrons need bigger chairs, so it’s the electrons we most have to worry about. As a result of the exclusion principle the electrons exert a degeneracy pressure to keep all the other electrons out of their chair. You can therefore treat the electrons as a simple gas with a pressure due to the exclusion principle.  Such a star is sometimes called a degenerate star, but it is more commonly known as a white dwarf.

A plot of the radius of a white dwarf vs its mass in a simple model.

So to model such a star, we just need gravity and the equation of state (a description of the pressure) for our electron gas.  What we find is that the more mass a white dwarf has, the smaller its radius. This is exactly what we would expect, since more mass means stronger gravity, which means all those degenerate electrons are squeezed harder. But when the mass reaches about 1.4 solar masses, the electrons are squeezed so tightly that the resulting radius is zero. At this limit the electron pressure can’t withstand gravity. This limit is known as the Chandrasekhar limit (named after its discoverer). Above this limit and the electrons are squeezed into the protons in the ions, leaving a massive ball of neutrons known as a neutron star. Actually, it isn’t quite that simple, but you get the idea.

Since a white dwarf can’t produce heat through fusion or anything else, once it forms it gradually cools.  Since white dwarfs are small they don’t have a lot of surface area, so they cool rather slowly, but given enough time such a star would cool to the background temperature of the universe, becoming a cold and dark frozen star.

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What’s a superhump?  And what has it got to do with astronomy?

It all has to do with binary stars.  Specifically a binary system where one of the stars (usually a white dwarf) is capturing material from the other, as depicted in the image below.  These binary systems are often called cataclysmic variables, or cataclysmic binaries, because the light and energy they give off can vary significantly.

As the white dwarf captures material from its companion, the material forms an accretion disk around the white dwarf.  Occasionally a hotspot can form in the accretion disk, and this causes a gradual variation in the brightness of the binary known as a hump.  As the pair orbits, the accretion disk (and the hot spot) becomes blocked by the companion, and then visible again from our vantage point, thus the “hump” in brightness.

Sometimes the accretion disk can have an even bigger hotspot known as an outburst.  In some of these outbursts the accretion disk can become unstable and basically rip apart.  Again there is a variation in brightness, but the variation is much larger, and is called a superhump.  What’s interesting about superhumps is that the variation is almost(but not quite) equal to the orbital period of the binary system.  If the brightness variation was simply due to the (now ripped apart) accretion disk moving in front of and behind the companion star, you would think the superhump would follow the orbital period.  But since it is a bit off, what gives?

A recent paper in Acta Astronomica looks at this issue.  The working model first proposed in the 1990s was that the difference is due to tidal effects.  Basically a resonance would build between the accretion disk and companion star causing the accretion disk to shift over time (apisidal precession).  During an outburst this causes a bulge in the accretion material that leads to the variation in brightness.

But this alone can’t account for the brightness of the superhumps, as the paper points out.  It seems there is also an interaction between the accretion of new material and the outburst of the accretion disk.  As the bulge forms in the accretion disk, this affects the influx of new material, and the combined effects lead to the variation in brightness we see from the system.

It is clear there are some complex interactions happening under a rather silly name.

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