The expanding remnant of SN 1987A, a Type II-P supernova in the Large Magellanic Cloud. NASA image.
A Type II supernova (plural: supernovae or supernovas) results from the rapid collapse and violent explosion of a massive star. A star must have at least 8 times, but no more than 40 to 50 times, the mass of the Sun (M☉) to undergo this type of explosion. Type II supernovae are distinguished from other types of supernovae by the presence of hydrogen in their spectra. They are usually observed in the spiral arms of galaxies and in H II regions, but not in elliptical galaxies;
those are generally composed of older low-mass stars, with few of the
young highly massive stars necessary to cause a supernova.
Stars generate energy by the nuclear fusion of elements. Unlike the Sun, massive stars possess the mass needed to fuse elements that have an atomic mass greater than hydrogen and helium, albeit at increasingly higher temperatures and pressures, causing increasingly shorter stellar life spans. The degeneracy pressure of electrons and the energy generated by these fusion reactions
are sufficient to counter the force of gravity and prevent the star
from collapsing, maintaining stellar equilibrium. The star fuses
increasingly higher mass elements, starting with hydrogen and then helium, progressing up through the periodic table until a core of iron and nickel
is produced. Fusion of iron or nickel produces no net energy output, so
no further fusion can take place, leaving the nickel–iron core inert.
Due to the lack of energy output creating outward thermal pressure, the
core contracts due to gravity until the overlying weight of the star can
be supported largely by electron degeneracy pressure.
When the compacted mass of the inert core exceeds the Chandrasekhar limit of about 1.4 M☉, electron degeneracy is no longer sufficient to counter the gravitational compression. A cataclysmic implosion
of the core takes place within seconds. Without the support of the
now-imploded inner core, the outer core collapses inwards under gravity
and reaches a velocity of up to 23% of the speed of light and the sudden compression increases the temperature of the inner core to up to 100 billion kelvins. Neutrons and neutrinos are formed via reversed beta-decay, releasing about 1046 joules (100 foe) in a ten-second burst. Also, the collapse of the inner core is halted by neutron degeneracy, causing the implosion to rebound and bounce outward. The energy of this expanding shock wave
is sufficient to disrupt the overlying stellar material and accelerate
it to escape velocity, forming a supernova explosion. The shock wave and
extremely high temperature and pressure rapidly dissipate but are
present for long enough to allow for a brief period during which the
production of elements heavier than iron occurs. Depending on initial size of the star, the remnants of the core form a neutron star or a black hole. Because of the underlying mechanism, the resulting supernova is also described as a core-collapse supernova.
There exist several categories of Type II supernova explosions, which are categorized based on the resulting light curve—a graph of luminosity versus time—following the explosion. Type II-L supernovae show a steady (linear)
decline of the light curve following the explosion, whereas Type II-P
display a period of slower decline (a plateau) in their light curve
followed by a normal decay. Type Ib and Ic supernovae
are a type of core-collapse supernova for a massive star that has shed
its outer envelope of hydrogen and (for Type Ic) helium. As a result,
they appear to be lacking in these elements.
Formation
The onion-like layers of a massive, evolved star just before core collapse. (Not to scale.)
Stars far more massive than the sun evolve in more complex ways. In the core of the star, hydrogen is fused into helium, releasing thermal energy that heats the sun's core and provides outward pressure that supports the sun's layers against collapse in a process known as stellar or hydrostatic equilibrium.
The helium produced in the core accumulates there since temperatures in
the core are not yet high enough to cause it to fuse. Eventually, as
the hydrogen at the core is exhausted, fusion starts to slow down, and gravity
causes the core to contract. This contraction raises the temperature
high enough to initiate a shorter phase of helium fusion, which accounts
for less than 10% of the star's total lifetime. In stars with fewer
than eight solar masses, the carbon produced by helium fusion does not fuse, and the star gradually cools to become a white dwarf. White dwarf stars, if they have a near companion, may then become Type Ia supernovae.
A much larger star, however, is massive enough to create
temperatures and pressures needed to cause the carbon in the core to
begin to fuse when the star contracts at the end of the helium-burning
stage. The cores of these massive stars become layered like onions as
progressively heavier atomic nuclei build up at the center, with an
outermost layer of hydrogen gas, surrounding a layer of hydrogen fusing
into helium, surrounding a layer of helium fusing into carbon via the triple-alpha process,
surrounding layers that fuse to progressively heavier elements. As a
star this massive evolves, it undergoes repeated stages where fusion in
the core stops, and the core collapses until the pressure and
temperature are sufficient to begin the next stage of fusion, reigniting
to halt collapse.
Core-burning nuclear fusion stages for a 25-solar mass star Process Main fuel Main products 25 M☉ star Temperature
(K)Density
(g/cm3)Duration hydrogen burning hydrogen helium 7×107 10 107 years triple-alpha process helium carbon, oxygen 2×108 2000 106 years carbon burning process carbon Ne, Na, Mg, Al 8×108 106 1000 years neon burning process neon O, Mg 1.6×109 107 3 years oxygen burning process oxygen Si, S, Ar, Ca 1.8×109 107 0.3 years silicon burning process silicon nickel (decays into iron) 2.5×109 108 5 days
Core collapse
The factor limiting this process is the amount of energy that is released through fusion, which is dependent on the binding energy
that holds together these atomic nuclei. Each additional step produces
progressively heavier nuclei, which release progressively less energy
when fusing. In addition, from carbon-burning onwards, energy loss via neutrino production becomes significant, leading to a higher rate of reaction than would otherwise take place. This continues until nickel-56 is produced, which decays radioactively into cobalt-56 and then iron-56 over the course of a few months. As iron and nickel have the highest binding energy per nucleon of all the elements, energy cannot be produced at the core by fusion, and a nickel-iron core grows.
This core is under huge gravitational pressure. As there is no fusion
to further raise the star's temperature to support it against collapse,
it is supported only by degeneracy pressure of electrons. In this state, matter is so dense that further compaction would require electrons to occupy the same energy states. However, this is forbidden for identical fermion particles, such as the electron – a phenomenon called the Pauli exclusion principle.
When the core's mass exceeds the Chandrasekhar limit of about 1.4 M☉, degeneracy pressure can no longer support it, and catastrophic collapse ensues. The outer part of the core reaches velocities of up to 70000 km/s (23% of the speed of light) as it collapses toward the center of the star. The rapidly shrinking core heats up, producing high-energy gamma rays that decompose iron nuclei into helium nuclei and free neutrons via photodisintegration. As the core's density increases, it becomes energetically favorable for electrons and protons to merge via inverse beta decay, producing neutrons and elementary particles called neutrinos.
Because neutrinos rarely interact with normal matter, they can escape
from the core, carrying away energy and further accelerating the
collapse, which proceeds over a timescale of milliseconds. As the core
detaches from the outer layers of the star, some of these neutrinos are
absorbed by the star's outer layers, beginning the supernova explosion.
For Type II supernovae, the collapse is eventually halted by
short-range repulsive neutron-neutron interactions, mediated by the strong force, as well as by degeneracy pressure
of neutrons, at a density comparable to that of an atomic nucleus. When
the collapse stops, the infalling matter rebounds, producing a shock wave
that propagates outward. The energy from this shock dissociates heavy
elements within the core. This reduces the energy of the shock, which
can stall the explosion within the outer core.
The core collapse phase is so dense and energetic that only
neutrinos are able to escape. As the protons and electrons combine to
form neutrons by means of electron capture,
an electron neutrino is produced. In a typical Type II supernova, the
newly formed neutron core has an initial temperature of about 100
billion kelvins, 104
times the temperature of the Sun's core. Much of this thermal energy
must be shed for a stable neutron star to form, otherwise the neutrons
would "boil away". This is accomplished by a further release of
neutrinos. These 'thermal' neutrinos form as neutrino-antineutrino pairs of all flavors, and total several times the number of electron-capture neutrinos. The two neutrino production mechanisms convert the gravitational potential energy of the collapse into a ten-second neutrino burst, releasing about 1046 joules (100 foe).
Through a process that is not clearly understood, about 1%, or 1044 joules
(1 foe), of the energy released (in the form of neutrinos) is
reabsorbed by the stalled shock, producing the supernova explosion. Neutrinos generated by a supernova were observed in the case of Supernova 1987A, leading astrophysicists to conclude that the core collapse picture is basically correct. The water-based Kamiokande II and IMB instruments detected antineutrinos of thermal origin, while the gallium-71-based Baksan instrument detected neutrinos (lepton number = 1) of either thermal or electron-capture origin.
Within
a massive, evolved star (a) the onion-layered shells of elements
undergo fusion, forming a nickel-iron core (b) that reaches
Chandrasekhar-mass and starts to collapse. The inner part of the core is
compressed into neutrons (c), causing infalling material to bounce (d)
and form an outward-propagating shock front (red). The shock starts to
stall (e), but it is re-invigorated by neutrino interaction. The
surrounding material is blasted away (f), leaving only a degenerate
remnant.
When the progenitor star is below about 20 M☉ –
depending on the strength of the explosion and the amount of material
that falls back – the degenerate remnant of a core collapse is a neutron star. Above this mass, the remnant collapses to form a black hole. The theoretical limiting mass for this type of core collapse scenario is about 40–50 M☉. Above that mass, a star is believed to collapse directly into a black hole without forming a supernova explosion, although uncertainties in models of supernova collapse make calculation of these limits uncertain.
Theoretical models
The Standard Model of particle physics is a theory which describes three of the four known fundamental interactions between the elementary particles that make up all matter.
This theory allows predictions to be made about how particles will
interact under many conditions. The energy per particle in a supernova
is typically 1–150 picojoules (tens to hundreds of MeV).[18][failed verification]
The per-particle energy involved in a supernova is small enough that
the predictions gained from the Standard Model of particle physics are
likely to be basically correct. But the high densities may require
corrections to the Standard Model. In particular, Earth-based particle accelerators can produce particle interactions which are of much higher energy than are found in supernovae,
but these experiments involve individual particles interacting with
individual particles, and it is likely that the high densities within
the supernova will produce novel effects. The interactions between
neutrinos and the other particles in the supernova take place with the weak nuclear force, which is believed to be well understood. However, the interactions between the protons and neutrons involve the strong nuclear force, which is much less well understood.
The major unsolved problem with Type II supernovae is that it is not understood how the burst of neutrinos
transfers its energy to the rest of the star producing the shock wave
which causes the star to explode. From the above discussion, only one
percent of the energy needs to be transferred to produce an explosion,
but explaining how that one percent of transfer occurs has proven
extremely difficult, even though the particle interactions involved are
believed to be well understood. In the 1990s, one model for doing this
involved convective overturn, which suggests that convection, either from neutrinos
from below, or infalling matter from above, completes the process of
destroying the progenitor star. Heavier elements than iron are formed
during this explosion by neutron capture, and from the pressure of the
neutrinos pressing into the boundary of the "neutrinosphere", seeding
the surrounding space with a cloud of gas and dust which is richer in
heavy elements than the material from which the star originally formed.
Neutrino physics, which is modeled by the Standard Model, is crucial to the understanding of this process. The other crucial area of investigation is the hydrodynamics of the plasma that makes up the dying star; how it behaves during the core collapse determines when and how the shockwave forms and when and how it stalls and is reenergized.
In fact, some theoretical models incorporate a hydrodynamical
instability in the stalled shock known as the "Standing Accretion Shock
Instability" (SASI). This instability comes about as a consequence of
non-spherical perturbations oscillating the stalled shock thereby
deforming it. The SASI is often used in tandem with neutrino theories in
computer simulations for re-energizing the stalled shock.
Computer models
have been very successful at calculating the behavior of Type II
supernovae when the shock has been formed. By ignoring the first second
of the explosion, and assuming that an explosion is started, astrophysicists have been able to make detailed predictions about the elements produced by the supernova and of the expected light curve from the supernova.
Light curves for Type II-L and Type II-P supernovae
This
graph of the luminosity as a function of time shows the characteristic
shapes of the light curves for a Type II-L and II-P supernova.
When the spectrum of a Type II supernova is examined, it normally displays Balmer absorption lines – reduced flux at the characteristic frequencies where hydrogen atoms absorb energy. The presence of these lines is used to distinguish this category of supernova from a Type I supernova.
When the luminosity of a Type II supernova is plotted over a
period of time, it shows a characteristic rise to a peak brightness
followed by a decline. These light curves have an average decay rate of
0.008 magnitudes
per day; much lower than the decay rate for Type Ia supernovae. Type II
is subdivided into two classes, depending on the shape of the light
curve. The light curve for a Type II-L supernova shows a steady (linear)
decline following the peak brightness. By contrast, the light curve of a
Type II-P supernova has a distinctive flat stretch (called a plateau)
during the decline; representing a period where the luminosity decays
at a slower rate. The net luminosity decay rate is lower, at
0.0075 magnitudes per day for Type II-P, compared to 0.012 magnitudes
per day for Type II-L.
The difference in the shape of the light curves is believed to be
caused, in the case of Type II-L supernovae, by the expulsion of most
of the hydrogen envelope of the progenitor star. The plateau phase in Type II-P supernovae is due to a change in the opacity of the exterior layer. The shock wave ionizes
the hydrogen in the outer envelope – stripping the electron from the
hydrogen atom – resulting in a significant increase in the opacity.
This prevents photons from the inner parts of the explosion from
escaping. When the hydrogen cools sufficiently to recombine, the outer
layer becomes transparent.
Type IIn supernovae
The
"n" denotes narrow, which indicates the presence of narrow or
intermediate width hydrogen emission lines in the spectra. In the
intermediate width case, the ejecta from the explosion may be
interacting strongly with gas around the star – the circumstellar
medium.
The estimated circumstellar density required to explain the
observational properties is much higher than that expected from the
standard stellar evolution theory.
It is generally assumed that the high circumstellar density is due to
the high mass-loss rates of the Type IIn progenitors. The estimated
mass-loss rates are typically higher than 10−3 M☉ per year. There are indications that they originate as stars similar to Luminous blue variables with large mass losses before exploding. SN 1998S and SN 2005gl are examples of Type IIn supernovae; SN 2006gy, an extremely energetic supernova, may be another example.
Type IIb supernovae
A
Type IIb supernova has a weak hydrogen line in its initial spectrum,
which is why it is classified as a Type II. However, later on the H
emission becomes undetectable, and there is also a second peak in the
light curve that has a spectrum which more closely resembles a Type Ib supernova.
The progenitor could have been a massive star that expelled most of its
outer layers, or one which lost most of its hydrogen envelope due to
interactions with a companion in a binary system, leaving behind the
core that consisted almost entirely of helium. As the ejecta of a Type IIb expands, the hydrogen layer quickly becomes more transparent and reveals the deeper layers.
The classic example of a Type IIb supernova is SN 1993J, while another example is Cassiopeia A. The IIb class was first introduced (as a theoretical concept) by Woosley et al. in 1987, and the class was soon applied to SN 1987K and SN 1993J.
