We apologize for the absence, as we have been extremely and unexpectedly busy.

It’s Christmas time folks!  And what do we have at the top of the tree (if you celebrate)?  A star, right?  The perfect reminder to bring us back to our lovely discussion on stars!  So, we have discussed the life of a star.  Now things get interesting.  We have reached the end of a star’s lonely existence.  Yet another difference between a star’s end and a human’s end is that in some cases a star can have one of the most awesome events in the universe: a supernova.  Summarized, it is an explosion of a star.  But in detail there are multiple types, all fascinating.

Take that and put it on your Christmas tree! Okay, maybe you don’t celebrate or it doesn’t seem that festive, but it’s still awesome. From:

Each type has a progenitor, which makes spectra important as they show the amount of elements present and rate of fusion.  For example, during an SN enough energy is released such that higher than average fusion can take place.  So, if you ever wondered where iron or even heavier than iron elements came from, thank SNe.  Other ways to determine the type follow similar reasoning, including light curves, energy, and progenitor events.  The main summary of the types of SNe is the popular Minkowski-Zwicky scheme or classification. We will now begin splitting up these wonderful things, despite the fact that SNe involve fusion.  There are two types: I and II.  A major way to characterize them is by spectra.  Type I SNe lack Balmer hydrogen emission lines, while type II SNe show Balmer hydrogen emission lines in their spectra.


The Type Ia (thermonuclear) SN is the most different.  It consists of a binary pair of two stars, one being a white dwarf.  As stated previously in the WD, PN, Type Ia post , they go through an uncontrolled thermonuclear explosion as the white dwarf becomes unstable upon reaching the Chandrasekhar limit.  The mechanism of explosion is important to note as the other three main types explode from gravitational collapse.

The other three main types are known as core-collapse SNe, which we explained at the end of our post on high-mass stellar evolution.  As the name indicates they collapse from the sheer amount of matter in one space, creating the appropriately named implosion-explosion event.  Type Ib and Ic are thought to mostly come from WR stars, so they appeared to be very similar.  But they were differentiated by spectra, where the Type Ib has shed its hydrogen shell before collapsing, while the Type Ic has shed its helium shell too. (Despite being caused by totally different mechanisms, both Type Ia and Type Ic lack hydrogen and helium lines; they are differentiated by the strong silicon lines in Type Ia.) As we said last time, the shedding of layers before explosion makes them sometimes named stripped-core collapse SNe.  Type II SNe are massive stars, of generally about 8 to 12 solar masses that have enough mass to collapse and then either continue collapsing or explode.

They also differ by location.  Type Ia are located in most areas as they only require a white dwarf and a binary system, and this is a more likely outcome for most stars.  Type Ib, Ic, and II all are generally in the galactic spiral arms in areas of relatively recent star formation since they require more massive, population I stars, which have to be massive enough to explode.

Another difference involves energy and light curves.  Type Ia SNe have the highest and most consistent peak; the rest depend on mass.  Also, out of the SNe, the only one with the possibility for a plateau or non-linear falling luminosity is the type II SN.  Visually the light curves are like so (click to enlarge):

From Wikipedia Supernova Light Curves

Next is the Type II SN and its core-collapse.  As the core collapses, protons and electrons are pushed together, forming neutrons and neutrinos.  This accounts for the high neutrino signatures that come from these types.  In fact, the relatively recent issue of faster than light neutrinos involved this very type of SN!  The neutrinos have such high densities that they develop radiation pressure, but the star essentially collapses in on itself, rebounds, and creates a shock wave.  These effects are named neutrino outburst and rebound shock.  Type II SNe can also be split up into types P, L, n, and b.  Type II-P and II-L stand for plateau and linear, which are the shapes of the light curve. The plateau type forms from hydrogen ionizing after the explosion, increasing luminosity, while the linear type fully expels its hydrogen layer.  Type IIn SNe are detected by narrow hydrogen emission lines in spectra, sometimes thought to be caused by LBVs.  Type IIb SNe show weak hydrogen lines initially, but the lines later become untraceable and it starts to resemble the spectra of a Type Ib SN.  The famous example of this is Cassiopeia A.  Another very high mass explosion is the hypernova, or a pair instability SN.  The pair instability occurs when atoms collide with gamma rays to form electrons and positrons, which creates a pressure to slightly reduce the core’s overall pressure.  This leads to a large-scale collapse.

For all the supernovae remnants (SNRs) can form.  For type Ia SNe they will always be large nebulae, but the other types can leave massive objects too.  These include a neutron star or a black hole.  The cause for the neutron star is that neutronization or neutrino outburst from neutron degeneracy pressure, analogous to the electron degeneracy pressure holding up a white dwarf, stops full collapse.  Black holes, on the other hand, collapse to a point.  Both objects have very odd and interesting properties.

To finish off we aren’t fooling that there are also supernova impostors, originally thought to be a type V SN.  These are also possibly LBV eruptions, but they don’t seem to be full supernovae.  As stated LBVs are unstable, so they can throw off large amounts of mass and heat that up.  They are recent but important events that are still being researched.



The easiest summary of supernovae is that they are characterized by many ways.  But they have great importance.  They explain how we could have any jewelry (in case you wanted to get some for a holiday present).  Even greater, they allow for a better understanding of the universe.  This is by giving ways to calibrate distance measures and  understand properties of stars.  Just to the point, they’re awesome.  It’s not Thanksgiving, but we should certainly thank SNe anyway.


References and further reading:

SNe in general:

Type Ib+Ic SNe:

Type II SNe:

SN imposter:

High-Mass Stellar Evolution

We’ve spent the past couple weeks discussing low-mass stellar evolution and its results, but now it’s time to move onto bigger and better things. That’s right, now we’re going to explore high-mass stellar evolution, otherwise known as “the stars that do go BANG”.

First, we should clarify that by “high-mass stars“, we mean those with masses from approximately 8 M_sun to 100+ M_sun. The upper mass limit for a star is not definitively known, but it is approximated by the Eddington Limit, where the radiation pressure outwards is so strong that it overcomes inward gravity and blows the star apart.


High-mass stars form just like low-mass stars — only with, well, more mass. Also like low mass stars, they start their lives on the Main Sequence, fusing hydrogen to helium. However, high-mass stars have much shorter Main Sequence lifetimes than low-mass stars.  This is because even though they have more mass, that additional mass results in an increased rate of fusion which burns through their fuel faster. Thus, high-mass stars quickly leave the Main Sequence and become RSGs/BSGs, LBVs, WR stars, and SNe (yes, we really do love our acronyms in astronomy).

High-mass RSG structure

The inner structure of a high-mass red supergiant (Credit: New Mexico State University)

RSG and BSG stand for “red supergiant” and “blue supergiant“, respectively. Just like low-mass stars, high-mass stars also migrate to the right side of the H-R diagram when a helium core builds up and hydrogen shell-burning starts. The increased radiation pressure from faster fusion causes the star’s outer layers to expand and cool. So, we have a RSG. When the helium core ignites (this time without an explosive flash, since the core is not degenerate), the star heads back towards the “blue” part of the H-R diagram. So, we have a BSG. Then, when the star can no longer fuse helium in its core, it will become a RSG again; when the carbon core ignites, it will revert to a BSG; and so on. A high-mass star will make many treks across the H-R diagram from “blue” to “red” and back as successively heavier elements fuse in its core, with each journey producing less energy and lasting for a shorter period of time, until it reaches an iron core…

LBV Eta Carinae

Luminous Blue Variable Eta Carinae (Credit: CXO/NASA)

LBVs are Luminous Blue Variables (also called Hubble-Sandage variables, or S Doradus variables), which are basically exactly what they sound like — they are massive, luminous stars that are blue in color and vary in luminosity over time.  They spend most of their time in a quiescent state, varying by a magnitude or two over a period of years.   But occasionally they undergo violent outbursts, which can be so bright that they have repeatedly been mistaken for supernovae. Quiescent LBVs generally appear as B-type supergiants that shed mass at high rates, but during an eruption, they become cooler, become  redder and eject huge amounts of mass. It is thought that really massive stars pass through a short LBV phase on their way to becoming WR stars, but no one knows for sure.

WR star in NGC 2359

WR star in NGC 2359 (Credit: P. Berlind and P. Challis of the Harvard-Smithsonian Center for Astrophysics)

WR stars are Wolf-Rayet stars, named for astronomers Charles Wolf and Georges Rayet, who discovered these fascinating stars. These massive and incredibly luminous stars are primarily noted for their powerful stellar winds, which cause extremely high rates of mass loss. Unlike normal stars, WR stars have prominent lines of helium in their spectra, along with nitrogen (WN), carbon (WC), or rarely oxygen (WO) (there is logic here since these elements are in the CNO cycle). WN stars can be further divided into “late” classes (L) and “early” (E) – WNL stars have hydrogen emission lines in their spectra, while WNE stars do not. Like the LBV stage, high-mass stars do not spend a long time in the WR stage, so these weird stars are quite rare.

SNe of course stands for those wondrous exploding stars, supernovae! Massive stars are thought to be the causes of every type of supernovae except Type Ia (if you will recall, those are caused by exploding white dwarfs).  These include Type Ib and Ic, and all the subtypes of Type II.

The precise stages that a high-mass star will pass through depend on its mass. There is significant variability in the exact masses and stages from source to source, so we have avoided trying to compile everything into one table because it would just end up as a mess. In general, stars on the low end of the “high-mass” spectrum tend to turn into RSGs or BSGs, then going supernova. More massive high-mass stars generally evolve into LBVs, then WN stars, then WC stars (or very occasionally, WO stars), before finally exploding as supernovae.


While all these changes are taking place in the exterior of the star, the stellar interior is constantly fusing heavier and heavier elements. As we mentioned in the post about nuclear reactions, massive stars above 8 M_sun can fuse elements all the way up to iron (technically up to cobalt and nickel, but these isotopes are radioactive and decay into iron). Any fusion past the so-called “iron peak” consumes energy instead of releasing it, which undermines exactly what the star has been trying to accomplish throughout its entire life.

When the core becomes hot and dense enough to fuse iron, photodisintegration (the splitting of elements by, you guessed it, photons or light) of heavy elements into light elements consumes energy and recently-formed neutrinos carry energy away, causing the core to collapse. The core will eventually reach such a high density that it refuses to be compressed any further and “bounces”, sending a shock wave outwards into the rest of the collapsing star. The shock wave stalls, but then accelerates again due to the neutrinos that are forcing their way outwards as well (neutrinos typically don’t like to interact with regular matter, but since there are so many of them here, they have quite an appreciable effect) — this is a core-collapse supernova. The exact mechanism of collapse and rebound for a supernova is not fully known, but overall this is what’s expected to occur…and hey, we are talking about giant massive stars, give some credit that our good astronomers could even find this much.  When stars have their outer envelopes shed, they are instead called stripped core-collapse supernovae (with this massive stars encompass all the non-Type Ia SNe).

M1 Crab Nebula

The famous M1 Crab Nebula, created by a core-collapse Type II supernova (Credit: HST)

After exploding, the core either forms a giant supernova remnant (SNR), becomes a neutron star or has so much mass that it collapses further into a black hole.  We shall save these interesting objects for another post, as this post would be intolerably long if we tried to discuss those here as well. (We will also cover all the different types of supernovae in a further post, since they are so important that they deserve a post of their own.)


TL;DR — High-mass stars start off their lives very much like low-mass stars.  While low-mass stars will become a red supergiant, eject PNe, and turn into WDs,  high-mass stars cycle between red and blue supergiants or turn into Luminous Blue Variables and Wolf-Rayet stars. In the end, high-mass stars will explode in (sometimes stripped) core-collapse supernovae and produce strange stellar remnants.


Links and sources for further reading

General High-Mass Evolution:

Luminous Blue Variables:

Wolf-Rayet stars:

Type II SNe:

Carroll and Ostlie, An Introduction to Modern Astrophysics, 2nd edition, p. 518 to 534

White Dwarfs, Planetary Nebulae and Type Ia Supernovae

In our last post about low-mass stellar evolution, we left off at the AGB.  So, what comes next?  We said of white dwarfs (WDs) and planetary nebulae (PNe), but these can get somewhat detailed, so they deserve a post of their own.  Especially since WDs have a very interesting relationship to a certain event we call a supernova (SN; we adore our acronyms).

After a star gets to the pulsating red giant phase, it eventually gets to the point where it doesn’t have the energy to continue fusion, usually by the C/N/O elements.  When this happens there is insufficient radiation pressure outwards to repel the gravity pushing the star inwards, making it condense (sadly we can’t condense all this astronomy information).  With that, the outer layers of the star shed off to form a PN.  Some may ask why they are called planetary nebulae, since they in fact do not have any planets at all.  It comes from the fact that early astronomers couldn’t see them clearly, and with PNe being generally green-blue, they looked similar to the planet Uranus.

Most PNe are indicators of the later stages of a low-mass star.  Therefore, they are generally found with Population I stars, younger stars closer to the galactic plane or center.  Like other nebulae, they also look awesome and are among the most popular targets for astrophotography.  An interesting fact is that they don’t only form as a spherical shell, but they can come in many shapes, such as a bipolar shape, and also rings and ellipses.  In general, PNe are distinguished by shape, so another major question is how these shapes form.  Many of the different shapes are theorized to be from high-speed stellar winds, but the rings most likely come from an ejection of material that occurred at a different time from the main ejection.

As for their colorful nature, PNe have red emission lines from hydrogen and have green lines from doubly ionized oxygen, or OIII.  They are created from the star’s light ionizing atoms, ripping electrons away and then recombining.  The lines from a PN are also known as forbidden lines, not because they shouldn’t happen, but because in a lab we could normally only produce these spectra under high density.  Space, on the other hand, is an odd, low-density environment.

The many shapes, colors, and sizes of these PNs. They may not be planetary, but they are certainly stellar! Credit: Bruce Balick and HST from

There are also Fast-Low Ionizing Emission Regions (FLIERs).  These are strange red-colored regions of a PN that are exactly what their name says.  But sadly their origin is not fully understood.  Some say they seem like they should be moving outward, but they could be stationary or even moving inward.  So yes, this lovely universe of ours likes to seem to make little sense.

Saturn Nebula (NGC 7009)

You can see the FLIERs at the sides. So…perhaps you can see what’s going on here. Credit: David Darling


They can have multiple features.  They may seem like degenerates for being aged as they are, but really they are fascinating.  They can be categorized as many types of stars. They can go ahead and nova or even go SN (for fun, let’s just say that they are little things with quite the Napoleon complex).  Yes, my fellow astronomy-lovers, we have gotten to the white dwarf (WD).

WDs are the result of  a star collapsing in on itself due to lack of radiation pressure, disrupting hydrostatic equilibrium, and allowing gravity to push inwards.  They come from stars that have run out of fuel, and they’re very small, being around the mass of the Sun shrunk down to the size of the Earth.  They are made mostly of the C, N, O elements because by that point a medium-mass star can’t fuse much more.  The color of the WD comes from the fact that it’s quite hot. It may not seem like it, but WDs can actually be detected in X-ray better than in visible light.

So then, the next question would be why doesn’t the star just continue to collapse?  It doesn’t even have fusion occurring to oppose the collapse.  Well, WDs actually have stability. They fulfill hydrostatic equilibrium, due to the Pauli exclusion principle, which works independent of temperature.  The compaction of the WD produces such a high density that all the atoms are smashed together to the point where electrons start to collide.  The Pauli exclusion principle says that two particles cannot occupy the same space with the same energy and spin, since they are effectively identical.  So, what happens when you have a lot of electrons in a little space?  Some that are in the same state collide.  Blasphemy!  In actuality a pressure keeps these two electrons apart, which creates pressure that stops total collapse.  This state of matter, as mentioned in the light and matter post, is known as “degenerate matter.”

Similar picture with electrons

And to summarize this. There’s not much space, so we are under high pressure here! This can oppose even the strong force of gravity. Credit: Chandra X-ray

Next, how do we detect them?  Well, Friedrich Bessel did manage to spot the white dwarf companion of Sirius in the visible range through careful observation (we’ll call this one Sirius business).  WDs can (and often do) form in binary star systems, which is an especially important way to detect any sort of object that’s normally tough to see, because if we cannot see the target object, the companion is still affected by that object.

And yet another interesting part to the WD is the famous discovery of one Subrahmanyan Chandrasekhar.  To get to the point, he found that WDs actually have a mass limit.  His work was highly debated at the time, but now it is heavily accepted.  So just how large can a star be before its mass overcomes even the pressure trying to oppose gravity?  The answer is about 1.4 solar masses.

What’s interesting about this is that there are actually multiple types of WDs, because the Chandrasekhar Limit is an upper mass limit.  This means smaller stars (even red dwarfs) can still form WDs, which have differing amounts of the elements C, N, O, and Ne.  For an analogy of how massive a WD is imagine one teaspoon of WD matter.  This is about 16 tons.  That is, if we had a beach ball of WD matter, that would still be heavier than an ocean liner.

white dwarf types

And now the summary of WD types. Credit: David Darling

Ultimately these stars would form into black dwarfs, fading away into another boring death.  But does this always happen?  Well, since it’s the universe, of course not!  In fact, WDs not only have multiple types, but they can even pulsate.  This will be discussed in more detail later (see, we like to mimic stars and evolve this blog as we go along).  Aside from spectral differences from different elements they also can differ by the presence of magnetic fields.


But the mass limit has even more interesting implications.  See, WDs can be in binary systems with a red giant star with unstably held outer layers.  Those layers can be ripped off the companion star and accumulate around the WD in a process known as accretion.  Aside from our fancy names, accretion allows for material to get close enough to the WD such that fusion can occur.  This flash of very high amounts of fusion all at once is what we call a nova.  The name comes from the fact that when astronomers first saw these bright flashes, they thought new stars were being made.  But the even greater thing is that when a WD gains more than 1.4 solar masses, it produces one of the many magnificent events in the universe, an supernova.

Now, to be specific SNe can be split into two types, and they are further split up from there.  The type of SN produced by a WD is known as a Type Ia SN.  There is hot debate over what causes it, perhaps at a thermonuclear level!  Okay, not that hot, but one theory is the above mentioned accretion.  There is also what’s known as the double degenerate progenitor (a progenitor is a fancy word meaning “the thing that created it”), which states that a binary system of two WDs with a total mass greater than 1.4 solar masses collided and, exceeding the limit, exploded. But what is the exact cause of an SN?  Well, so much mass is gained in such a little space that runaway thermonuclear fusion occurs when it burns up all at once, releasing a massive amount of energy.

A nice summary of the SN’s creation. Credit: hyperphysics

Type Ia SNe are well-known for multiple reasons.  They are located almost everywhere, since WDs are just older low-mass main sequence stars.  They are thought to be able to fuse carbon/oxygen all the way up to iron/nickel, creating the heavy elements of the universe. The spectra is noted for a lack of Balmer hydrogen lines, and further distinguished by the presence of silicon in their spectra.  Another major note is that since there is a mass limit on WDs, the light curves of a Type Ia SN should always be the same. This means they are what’s called a standard candle; that is, they can be used to consistently measure what distance the SN was from earth.  The absolute magnitude for a Type Ia is approximately -19.5.  This is essential for understanding distances across the cosmos and can be used to calibrate Hubble’s Law.

Lastly, we will show how this is really, really powerful.  An initial comparison is that some note that this is enough to equal the entire Sun’s energy emitted over its whole lifetime.  But let’s continue to show another way.  We said earlier in the apparent/absolute magnitude post that we can relate brightness or luminosity to magnitudes.  Let’s apply this now.  Given the absolute magnitude of the Sun is +4.83 and a solar luminosity is about 4×1026 W we can now find it:


L of the SN/1 solar luminosity =100(4.83+19.5)/5

=5395106225.15 solar luminosities = 2×1036 W.

Now to put this relative.  The global energy consumption per year is about 15 terrawatts.  That is, we could run the entire world’s energy supply and then some  for oh about 1023 years.  I believe it is undebatable to say this is insane.  Now, others say that the energy output is about 2×1044 J.  So, take this all as you will.



Now to condense (perhaps into a star, or even a WD!) all this information.  After we have a red giant star it sheds its loosely held layers to form a PN.  This PN can have many different colors and shapes.  At the center the remaining core of mass condenses into a white dwarf, composed of degenerate matter.  This means it is so condensed electrons repel eachother, creating an electron degeneracy pressure opposing gravity.  Beyond that, WDs have a mass limit known as the Chandrasekhar limit, being 1.4 solar masses. WDs can also form in binary systems.  If the partner accretes, or accumulates, matter to the WD then the WD can quickly fuse it and form an extremely bright nova.  If either this goes past the Chandrasekhar limit OR if a WD combines with another WD and exceeds the limit then the result is a beautiful Type Ia SN.  Type Ia SNe are valuable for creating heavier than normal elements, being common, and being standard in peak luminosity.  This is useful for calculation of large distances which leads to a better understanding of what we like to call the universe.


Sources and further reading:

Planetary Nebulae:


White Dwarfs:

Type Ia SNe:

Carroll and Ostlie, An Introduction to Modern Astrophysics, 2nd edition (pg 446-474,  524-529, 557-578)