Potassium-40 (40K) is a radioactive isotope of potassium which has a long half-life of 1.251×109 years. It makes up 0.012% (120 ppm) of the total amount of potassium found in nature.
Potassium-40 is a rare example of an isotope that undergoes both types of beta decay. In about 89.28% of events it decays to calcium-40 (40Ca) with emission of a beta particle (β−, an electron) with a maximum energy of 1.31 MeV and an antineutrino. In about 10.72% of events it decays to argon-40 (40Ar) by electron capture (EC), with the emission of a neutrino and then a 1.460 MeV gamma ray. The radioactive decay of this particular isotope explains the large abundance of argon (nearly 1%) in the Earth's atmosphere, as well as prevalence of 40Ar over other isotopes. Very rarely (0.001% of events) it will decay to 40Ar by emitting a positron (β+) and a neutrino.
Potassium–argon dating
Decay scheme
Potassium-40 is especially important in potassium–argon (K–Ar) dating. Argon is a gas that does not ordinarily combine with other elements. So, when a mineral forms – whether from molten rock,
or from substances dissolved in water – it will be initially
argon-free, even if there is some argon in the liquid. However, if the
mineral contains any potassium, then decay of the 40K isotope
present will create fresh argon-40 that will remain locked up in the
mineral. Since the rate at which this conversion occurs is known, it is
possible to determine the elapsed time since the mineral formed by
measuring the ratio of 40K and 40Ar atoms contained in it.
The argon found in Earth's atmosphere is 99.6% 40Ar; whereas the argon in the Sun – and presumably in the primordial material that condensed into the planets – is mostly 36Ar, with less than 15% of 38Ar.
It follows that most of the terrestrial argon derives from potassium-40
that decayed into argon-40, which eventually escaped to the atmosphere.
Contribution to natural radioactivity
The evolution of Earth's mantle radiogenic heat flow over time: contribution from 40K in yellow.
The radioactive decay of 40K in the Earth's mantle ranks third, after 232Th and 238U, as the source of radiogenic heat.
The core also likely contains radiogenic sources, although how much is
uncertain. It has been proposed that significant core radioactivity
(1–2 TW) may be caused by high levels of U, Th, and K.
Potassium-40 is the largest source of natural radioactivity in
animals including humans. A 70 kg human body contains about 140 grams of
potassium, hence about 0.000117 × 140 = 0.0164 grams of 40K; whose decay produces about 4,300 disintegrations per second (becquerel) continuously throughout the life of the body.
Hubble Space Telescope images of HH 32 (left; top) and HH 34 (right; upper left) – colourful nebulae are typical of Herbig–Haro objects
Herbig–Haro (HH) objects are bright patches of nebulosity associated with newborn stars. They are formed when narrow jets of partially ionised gas
ejected by stars collide with nearby clouds of gas and dust at several
hundred kilometres per second. Herbig–Haro objects are commonly found in
star-forming regions, and several are often seen around a single star, aligned with its rotational axis. Most of them lie within about one parsec (3.26 light-years)
of the source, although some have been observed several parsecs away.
HH objects are transient phenomena that last around a few tens of
thousands of years. They can change visibly over timescales of a few
years as they move rapidly away from their parent star into the gas
clouds of interstellar space (the interstellar medium or ISM). Hubble Space Telescope
observations have revealed the complex evolution of HH objects over the
period of a few years, as parts of the nebula fade while others
brighten as they collide with the clumpy material of the interstellar
medium.
First observed in the late 19th century by Sherburne Wesley Burnham, Herbig–Haro objects were recognised as a distinct type of emission nebula in the 1940s. The first astronomers to study them in detail were George Herbig and Guillermo Haro, after whom they have been named. Herbig and Haro were working independently on studies of star formation
when they first analysed the objects, and recognised that they were a
by-product of the star formation process. Although HH objects are a
visible wavelength phenomena, many remain invisible at these wavelengths due to dust and gas, and can only be detected at infrared wavelengths. Such objects, when observed in near infrared, are called molecular hydrogen emission-line objects (MHOs).
Discovery and history of observations
The first HH object was observed in the late 19th century by Sherburne Wesley Burnham, when he observed the star T Tauri with the 36-inch (910 mm) refracting telescope at Lick Observatory and noted a small patch of nebulosity nearby. It was thought to be an emission nebula, later becoming known as Burnham's Nebula, and was not recognized as a distinct class of object. T Tauri was found to be a very young and variable star, and is the prototype of the class of similar objects known as T Tauri stars which have yet to reach a state of hydrostatic equilibrium between gravitational collapse and energy generation through nuclear fusion at their centres.
Fifty years after Burnham's discovery, several similar nebulae were
discovered with almost star-like appearance. Both Haro and Herbig made
independent observations of several of these objects in the Orion Nebula during the 1940s. Herbig also looked at Burnham's Nebula and found it displayed an unusual electromagnetic spectrum, with prominent emission lines of hydrogen, sulfur and oxygen. Haro found that all the objects of this type were invisible in infrared light.
Following their independent discoveries, Herbig and Haro met at an astronomy conference in Tucson, Arizona
in December 1949. Herbig had initially paid little attention to the
objects he had discovered, being primarily concerned with the nearby
stars, but on hearing Haro's findings he carried out more detailed
studies of them. The Soviet astronomer Viktor Ambartsumian
gave the objects their name (Herbig–Haro objects, normally shortened to
HH objects), and based on their occurrence near young stars (a few
hundred thousand years old), suggested they might represent an early
stage in the formation of T Tauri stars. Studies of the HH objects showed they were highly ionised, and early theorists speculated that they were reflection nebulae
containing low-luminosity hot stars deep inside. But the absence of
infrared radiation from the nebulae meant there could not be stars
within them, as these would have emitted abundant infrared light. In
1975 American astronomer R. D. Schwartz theorized that winds from T Tauri stars produce shocks in the ambient medium on encounter, resulting in generation of visible light. With the discovery of the first proto-stellarjet in HH 46/47, it became clear that HH objects are indeed shock-induced phenomena with shocks being driven by a collimated jet from protostars.
Formation
HH
objects are formed when accreted material is ejected by a protostar as
ionized gas along the star's axis of rotation, as exemplified by HH 34
(right).
Stars form by gravitational collapse of interstellar gas clouds. As the collapse increases the density, radiative energy loss decreases due to increased opacity.
This raises the temperature of the cloud which prevents further
collapse, and a hydrostatic equilibrium is established. Gas continues to
fall towards the core in a rotating disk. The core of this system is called a protostar. Some of the accreting material is ejected out along the star's axis of rotation in two jets of partially ionised gas (plasma).
The mechanism for producing these collimated bipolar jets is not
entirely understood, but it is believed that interaction between the
accretion disk and the stellar magnetic field accelerates some of the accreting material from within a few astronomical units
of the star away from the disk plane. At these distances the outflow is
divergent, fanning out at an angle in the range of 10−30°, but it
becomes increasingly collimated at distances of tens to hundreds of
astronomical units from the source, as its expansion is constrained. The jets also carry away the excess angular momentum
resulting from accretion of material onto the star, which would
otherwise cause the star to rotate too rapidly and disintegrate. When these jets collide with the interstellar medium, they give rise to the small patches of bright emission which comprise HH objects.
Electromagnetic emission from HH objects is caused when shock waves collide with the interstellar medium, creating what is called the "terminal working surfaces". The spectrum is continuous, but also has intense emission lines of neutral and ionized species. Spectroscopic observations of HH objects' doppler shifts indicate velocities of several hundred kilometers per second, but the emission lines in those spectra
are weaker than what would be expected from such high-speed collisions.
This suggests that some of the material they are colliding with is also
moving along the beam, although at a lower speed.
Spectroscopic observations of HH objects show they are moving away from
the source stars at speeds of several hundred kilometres per second. In recent years, the high optical resolution of the Hubble Space Telescope has revealed the proper motion (movement along the sky plane) of many HH objects in observations spaced several years apart.
As they move away from the parent star, HH objects evolve
significantly, varying in brightness on timescales of a few years.
Individual compact knots or clumps within an object may brighten and
fade or disappear entirely, while new knots have been seen to appear. This is most likely because of the precession of the jets and their pulsating, rather than steady, eruption from the parent stars.
Faster jets catch up with earlier slower jets, creating the so-called
"internal working surfaces", where streams of gas collide and generate
shock waves and consequently emissions.
The total mass being ejected by stars to form typical HH objects is estimated to be of the order of 10−8 to 10−6M☉ per year, a very small amount of material compared to the mass of the stars themselves but amounting to about 1–10% of the total mass accreted by the source stars in a year. Mass loss tends to decrease with increasing age of the source. The temperatures observed in HH objects are typically about 9,000–12,000 K, similar to those found in other ionized nebulae such as H II regions and planetary nebulae.
Densities, on the other hand, are higher than in other nebulae, ranging
from a few thousand to a few tens of thousands of particles per cm3, compared to a few thousand particles per cm3 in most H II regions and planetary nebulae.
Densities also decrease as the source evolves over time. HH objects consist mostly of hydrogen and helium, which account for about 75% and 24% of their mass respectively. Around 1% of the mass of HH objects is made up of heavier chemical elements, including oxygen, sulfur, nitrogen, iron, calcium and magnesium. Abundances of these elements, determined from emission lines of respective ions, are generally similar to their cosmic abundances. Many chemical compounds found in the surrounding interstellar medium, but not present in the source material, such as metal hydrides, are believed to have been produced by shock-induced chemical reactions.
Around 20–30% of the gas in HH objects is ionized near the source star,
but this proportion decreases at increasing distances. This implies the
material is ionized in the polar jet, and recombines as it moves away
from the star, rather than being ionized by later collisions. Shocking at the end of the jet can re-ionise some material, giving rise to bright "caps".
Numbers and distribution
HH
1/2 (top), HH 34 (left), and HH 47 (right) were numbered in order of
their discovery; it is estimated that there are up to 150,000 such
objects in the Milky Way.
HH objects are named approximately in order of their identification;
HH 1 and HH 2 being the earliest such objects to be identified. More than a thousand individual objects are now known. They are always present in star-forming H II regions, and are often found in large groups. They are typically observed near Bok globules (dark nebulae
which contain very young stars) and often emanate from them. Several HH
objects have been seen near a single energy source, forming a string of
objects along the line of the polar axis of the parent star.
The number of known HH objects has increased rapidly over the last few
years, but that is a very small proportion of the estimated up to
150,000 in the Milky Way, the vast majority of which are too far away to be resolved. Most HH objects lie within about one parsec of their parent star. Many, however, are seen several parsecs away.
HH 46/47 is located about 450 parsecs (1,500 light-years) away from the Sun and is powered by a class I protostarbinary.
The bipolar jet is slamming into the surrounding medium at a velocity
of 300 kilometers per second, producing two emission caps about 2.6
parsecs (8.5 light-years) apart. Jet outflow is accompanied by a 0.3
parsecs (0.98 light-years) long molecular gas outflow which is swept up
by the jet itself. Infrared studies by Spitzer Space Telescope have revealed a variety of chemical compounds in the molecular outflow, including water (ice), methanol, methane, carbon dioxide (dry ice) and various silicates. Located around 460 parsecs (1,500 light-years) away in the Orion nebula, HH 34
is produced by a highly collimated bipolar jet powered by a class I
protostar. Matter in the jet is moving at about 220 kilometers per
second. Two bright bow shocks,
separated by about 0.44 parsecs (1.4 light-years), are present on the
opposite sides of the source, followed by series of fainter ones at
larger distances, making the whole complex about 3 parsecs (9.8
light-years) long. The jet is surrounded by a 0.3 parsecs (0.98
light-years) long weak molecular outflow near the source.
Source stars
The stars from which HH jets are emitted are all very young stars, a
few tens of thousands to about a million years old. The youngest of
these are still protostars in the process of collecting from their
surrounding gases. Astronomers divide these stars into classes 0, I, II
and III, according to how much infrared radiation the stars emit.
A greater amount of infrared radiation implies a larger amount of
cooler material surrounding the star, which indicates it is still
coalescing. The numbering of the classes arises because class 0 objects
(the youngest) were not discovered until classes I, II and III had
already been defined.
Class 0 objects are only a few thousand years old; so young that
they are not yet undergoing nuclear fusion reactions at their centres.
Instead, they are powered only by the gravitational potential energy released as material falls onto them. They mostly contain molecular outflows with low velocities (less than a hundred kilometres per second) and weak emissions in the outflows.
Nuclear fusion has begun in the cores of Class I objects, but gas and
dust are still falling onto their surfaces from the surrounding nebula,
and most of their luminosity is accounted for by gravitational energy.
They are generally still shrouded in dense clouds of dust and gas, which
obscure all their visible light and as a result can only be observed at infrared and radio wavelengths. Outflows from this class are dominated by ionized species and velocities can range up to 400 kilometres per second.
The in-fall of gas and dust has largely finished in Class II objects
(Classical T Tauri stars), but they are still surrounded by disks of
dust and gas, and produce weak outflows of low luminosity. Class III objects (Weak-line T Tauri stars) have only trace remnants of their original accretion disk.
About 80% of the stars giving rise to HH objects are binary or
multiple systems (two or more stars orbiting each other), which is a
much higher proportion than that found for low mass stars on the main sequence.
This may indicate that binary systems are more likely to generate the
jets which give rise to HH objects, and evidence suggests the largest HH
outflows might be formed when multiple–star systems disintegrate.
It is thought that most stars originate from multiple star systems, but
that a sizable fraction of these systems are disrupted before their
stars reach the main sequence due to gravitational interactions with nearby stars and dense clouds of gas.
Infrared counterparts
HH
objects associated with very young stars or very massive protostars are
often hidden from view at optical wavelengths by the cloud of gas and
dust from which they form. The intervening material can diminish the visual magnitude
by factors of tens or even hundreds at optical wavelengths. Such deeply
embedded objects can only be observed at infrared or radio wavelengths, usually in the frequencies of hot molecular hydrogen or warm carbon monoxide emission.
In recent years, infrared images have revealed dozens of examples of
"infrared HH objects". Most look like bow waves (similar to the waves at
the head of a ship), and so are usually referred to as molecular "bow
shocks". The physics of infrared bow shocks can be understood in much
the same way as that of HH objects, since these objects are essentially
the same – supersonic shocks driven by collimated jets from the opposite poles of a protostar.
It is only the conditions in the jet and surrounding cloud that are
different, causing infrared emission from molecules rather than optical
emission from atoms and ions.
In 2009 the acronym "MHO", for Molecular Hydrogen emission-line Object,
was approved for such objects, detected in near infrared, by the International Astronomical Union
Working Group on Designations, and has been entered into their on-line
Reference Dictionary of Nomenclature of Celestial Objects. The MHO catalog contains over 2000 objects.
Seyfert galaxies are one of the two largest groups of active galaxies, along with quasars. They have quasar-like nuclei (very luminous, distant and bright sources of electromagnetic radiation) with very high surface brightnesses whose spectra reveal strong, high-ionisationemission lines, but unlike quasars, their host galaxies are clearly detectable.
Seyfert galaxies account for about 10% of all galaxies and are some of the most intensely studied objects in astronomy,
as they are thought to be powered by the same phenomena that occur in
quasars, although they are closer and less luminous than quasars. These
galaxies have supermassive black holes at their centers which are surrounded by accretion discs of in-falling material. The accretion discs are believed to be the source of the observed ultraviolet radiation. Ultraviolet emission and absorption lines provide the best diagnostics for the composition of the surrounding material.
Seen in visible light, most Seyfert galaxies look like normal spiral galaxies, but when studied under other wavelengths, it becomes clear that the luminosity of their cores is of comparable intensity to the luminosity of whole galaxies the size of the Milky Way.
Seyfert galaxies are named after Carl Seyfert, who first described this class in 1943.
Discovery
NGC 1068 (Messier 77), one of the first Seyfert galaxies classified
In 1926, Edwin Hubble looked at the emission lines of NGC 1068 and two other such "nebulae" and classified them as extragalactic objects. In 1943, Carl Keenan Seyfert
discovered more galaxies similar to NGC 1068 and reported that these
galaxies have very bright stellar-like nuclei that produce broad
emission lines. In 1944 Cygnus A was detected at 160 MHz, and detection was confirmed in 1948 when it was established that it was a discrete source. Its double radio structure became apparent with the use of interferometry. In the next few years, other radio sources such as supernova
remnants were discovered. By the end of the 1950s, more important
characteristics of Seyfert galaxies were discovered, including the fact
that their nuclei are extremely compact (< 100 pc, i.e.
"unresolved"), have high mass (≈109±1 solar masses), and the duration of peak nuclear emissions is relatively short (> 108 years).
NGC 5793 is a Seyfert galaxy located over 150 million light-years away in the constellation of Libra.
In the 1960s and 1970s, research to further understand the properties
of Seyfert galaxies was carried out. A few direct measurements of the
actual sizes of Seyfert nuclei were taken, and it was established that
the emission lines in NGC 1068 were produced in a region over a thousand
light years in diameter. Controversy existed over whether Seyfert redshifts were of cosmological origin.
Confirming estimates of the distance to Seyfert galaxies and their age
were limited since their nuclei vary in brightness over a time scale of a
few years; therefore arguments involving distance to such galaxies and
the constant speed of light cannot always be used to determine their
age.
In the same time period, research had been undertaken to survey,
identify and catalogue galaxies, including Seyferts. Beginning in 1967, Benjamin Markarian
published lists containing a few hundred galaxies distinguished by
their very strong ultraviolet emission, with measurements on the
position of some of them being improved in 1973 by other researchers. At the time, it was believed that 1% of spiral galaxies are Seyferts. By 1977, it was found that very few Seyfert galaxies are ellipticals, most of them being spiral or barred spiral galaxies. During the same time period, efforts have been made to gather spectrophotometric
data for Seyfert galaxies. It became obvious that not all spectra from
Seyfert galaxies look the same, so they have been subclassified
according to the characteristics of their emission spectra. A simple division into types I and II has been devised, with the classes depending on the relative width of their emission lines.
It has been later noticed that some Seyfert nuclei show intermediate
properties, resulting in their being further subclassified into types
1.2, 1.5, 1.8 and 1.9 (see Classification).
Early surveys for Seyfert galaxies were biased in counting only the
brightest representatives of this group. More recent surveys that count
galaxies with low-luminosity and obscured Seyfert nuclei suggest that
the Seyfert phenomenon is actually quite common, occurring in 16% ± 5%
of galaxies; indeed, several dozen galaxies exhibiting the Seyfert
phenomenon exist in the close vicinity (≈27 Mpc) of our own galaxy. Seyfert galaxies form a substantial fraction of the galaxies appearing in the Markarian catalog, a list of galaxies displaying an ultraviolet excess in their nuclei.
Characteristics
Optical and ultraviolet images of the black hole in the center of NGC 4151, a Seyfert Galaxy
An active galactic nucleus (AGN) is a compact region at the center of a galaxy that has a higher than normal luminosity over portions of the electromagnetic spectrum.
A galaxy having an active nucleus is called an active galaxy. Active
galactic nuclei are the most luminous sources of electromagnetic
radiation in the Universe, and their evolution puts constraints on
cosmological models. Depending on the type, their luminosity varies over
a timescale from a few hours to a few years. The two largest subclasses
of active galaxies are quasars and Seyfert galaxies, the main
difference between the two being the amount of radiation they emit. In a
typical Seyfert galaxy, the nuclear source emits at visible wavelengths
an amount of radiation comparable to that of the whole galaxy's
constituent stars, while in a quasar, the nuclear source is brighter
than the constituent stars by at least a factor of 100. Seyfert galaxies have extremely bright nuclei, with luminosities ranging between 108 and 1011 solar luminosities. Only about 5% of them are radio bright; their emissions are moderate in gamma rays and bright in X-rays. Their visible and infrared spectra shows very bright emission lines of hydrogen, helium, nitrogen, and oxygen. These emission lines exhibit strong Doppler broadening, which implies velocities from 500 to 4,000 km/s (310 to 2,490 mi/s), and are believed to originate near an accretion disc surrounding the central black hole.
A lower limit to the mass of the central black hole can be calculated using the Eddington luminosity. This limit arises because light exhibits radiation pressure. Assume that a black hole is surrounded by a disc of luminous gas.
Both the attractive gravitational force acting on electron-ion pairs in
the disc and the repulsive force exerted by radiation pressure follow
an inverse-square law. If the gravitational force exerted by the black
hole is less than the repulsive force due to radiation pressure, the
disc will be blown away by radiation pressure.
The
image shows a model of an active galactic nucleus. The central black
hole is surrounded by an accretion disc, which is surrounded by a torus.
The broad line region and narrow line emission region are shown, as
well as jets coming out of the nucleus.
Emissions
The
emission lines seen on the spectrum of a Seyfert galaxy may come from
the surface of the accretion disc itself, or may come from clouds of gas
illuminated by the central engine in an ionization cone. The exact
geometry of the emitting region is difficult to determine due to poor
resolution of the galactic center. However, each part of the accretion
disc has a different velocity relative to our line of sight, and the
faster the gas is rotating around the black hole, the broader the
emission line will be. Similarly, an illuminated disc wind also has a position-dependent velocity.
The narrow lines are believed to originate from the outer part of
the active galactic nucleus, where velocities are lower, while the
broad lines originate closer to the black hole. This is confirmed by the
fact that the narrow lines do not vary detectably, which implies that
the emitting region is large, contrary to the broad lines which can vary
on relatively short timescales. Reverberation mapping
is a technique which uses this variability to try to determine the
location and morphology of the emitting region. This technique measures
the structure and kinematics of the broad line emitting region by
observing the changes in the emitted lines as a response to changes in
the continuum. The use of reverberation mapping requires the assumption
that the continuum originates in a single central source.
For 35 AGN, reverberation mapping has been used to calculate the mass
of the central black holes and the size of the broad line regions.
In the few radio-loud Seyfert galaxies that have been observed, the radio emission is believed to represent synchrotron emission
from the jet. The infrared emission is due to radiation in other bands
being reprocessed by dust near the nucleus. The highest energy photons
are believed to be created by inverse Compton scattering by a high temperature corona near the black hole.
Classification
NGC 1097
is an example of a Seyfert galaxy. A supermassive black hole with a
mass of 100 million solar masses lies at the center of the galaxy. The
area around the black hole emits large amounts of radiation from the
matter falling into the black hole.
Seyferts were first classified as Type I or II, depending on the
emission lines shown by their spectra. The spectra of Type I Seyfert
galaxies show broad lines that include both allowed lines, like H I,
He I or He II and narrower forbidden lines, like O III. They show some
narrower allowed lines as well, but even these narrow lines are much
broader than the lines shown by normal galaxies. However, the spectra of
Type II Seyfert galaxies show only narrow lines, both permitted and
forbidden. Forbidden lines are spectral lines that occur due to electron transitions not normally allowed by the selection rules of quantum mechanics,
but that still have a small probability of spontaneously occurring. The
term "forbidden" is slightly misleading, as the electron transitions
causing them are not forbidden but highly improbable.
NGC 6300 is a Type II galaxy in the southern constellation of Ara.
In some cases, the spectra show both broad and narrow permitted
lines, which is why they are classified as an intermediate type between
Type I and Type II, such as Type 1.5 Seyfert. The spectra of some of
these galaxies have changed from Type 1.5 to Type II in a matter of a
few years. However, the characteristic broad Hα emission line has rarely, if ever, disappeared.
The origin of the differences between Type I and Type II Seyfert
galaxies is not known yet. There are a few cases where galaxies have
been identified as Type II only because the broad components of the
spectral lines have been very hard to detect. It is believed by some
that all Type II Seyferts are in fact Type I, where the broad components
of the lines are impossible to detect because of the angle we are at
with respect to the galaxy. Specifically, in Type I Seyfert galaxies, we
observe the central compact source more or less directly, therefore
sampling the high velocity clouds in the broad line emission region
moving around the supermassive black hole thought to be at the center of
the galaxy. By contrast, in Type II Seyfert galaxies, the active nuclei
are obscured and only the colder outer regions located further away
from the clouds' broad line emission region are seen. This theory is
known as the "Unification scheme" of Seyfert galaxies. However, it is not yet clear if this hypothesis can explain all the observed differences between the two types.
Type I Seyfert galaxies
NGC 6814 is a Seyfert galaxy with a highly variable source of X-ray radiation.
Type I Seyferts are very bright sources of ultraviolet light and X-rays
in addition to the visible light coming from their cores. They have two
sets of emission lines on their spectra: narrow lines with widths
(measured in velocity units) of several hundred km/s, and broad lines
with widths up to 104 km/s.
The broad lines originate above the accretion disc of the supermassive
black hole thought to power the galaxy, while the narrow lines occur
beyond the broad line region of the accretion disc. Both emissions are
caused by heavily ionised gas. The broad line emission arises in a
region 0.1-1 parsec across. The broad line emission region, RBLR,
can be estimated from the time delay corresponding to the time taken by
light to travel from the continuum source to the line-emitting gas.
Type II Seyfert galaxies
NGC 3081 is known as a Type II Seyfert galaxy, characterised by its dazzling nucleus.
Type II Seyfert galaxies have the characteristic bright core, as well as appearing bright when viewed at infrared wavelengths.
Their spectra contain narrow lines associated with forbidden
transitions, and broad lines associated with allowed strong dipole or
intercombination transitions. NGC 3147 is considered the best candidate to be a true Type II Seyfert galaxy. In some Type II Seyfert galaxies, analysis with a technique called spectro-polarimetry (spectroscopy of polarised light component) revealed obscured Type I regions. In the case of NGC 1068, nuclear light reflected off a dust cloud was measured, which led scientists to believe in the presence of an obscuring dust torus
around a bright continuum and broad emission line nucleus. When the
galaxy is viewed from the side, the nucleus is indirectly observed
through reflection by gas and dust above and below the torus. This reflection causes the polarisation.
In 1981, Donald Osterbrock
introduced the notations Type 1.5, 1.8 and 1.9, where the subclasses
are based on the optical appearance of the spectrum, with the
numerically larger subclasses having weaker broad-line components
relative to the narrow lines. For example, Type 1.9 only shows a broad component in the Hα line, and not in higher order Balmer lines. In Type 1.8, very weak broad lines can be detected in the Hβ
lines as well as Hα, even if they are very weak compared to the Hα. In
Type 1.5, the strength of the Hα and Hβ lines are comparable.
In addition to the Seyfert progression from Type I to Type II
(including Type 1.2 to Type 1.9), there are other types of galaxies that
are very similar to Seyferts or that can be considered as subclasses of
them. Very similar to Seyferts are the low-ionisation narrow-line
emission radio galaxies (LINER), discovered in 1980. These galaxies have
strong emission lines from weakly ionised or neutral atoms, while the
emission lines from strongly ionised atoms are relatively weak by
comparison. LINERs share a large amount of traits with low luminosity
Seyferts. In fact, when seen in visible light, the global
characteristics of their host galaxies are indistinguishable. Also, they
both show a broad line emission region, but the line emitting region in
LINERs has a lower density than in Seyferts. An example of such a galaxy is M104 in the Virgo constellation, also known as the Sombrero Galaxy. A galaxy that is both a LINER and a Type I Seyfert is NGC 7213, a galaxy that is relatively close compared to other AGNs.
Another very interesting subclass are the narrow line Type I galaxies
(NLSy1), which have been subject to extensive research in recent years.
They have much narrower lines than the broad lines from classic Type I
galaxies, steep hard and soft X-ray spectra and strong Fe[II] emission.
Their properties suggest that NLSy1 galaxies are young AGNs with high
accretion rates, suggesting a relatively small but growing central black
hole mass.
There are theories suggesting that NLSy1s are galaxies in an early
stage of evolution, and links between them and ultraluminous infrared
galaxies or Type II galaxies have been proposed.
Evolution
The majority of active galaxies are very distant and show large Doppler shifts. This suggests that active galaxies occurred in the early Universe and, due to cosmic expansion, are receding away from the Milky Way
at very high speeds. Quasars are the furthest active galaxies, some of
them being observed at distances 12 billion light years away. Seyfert
galaxies are much closer than quasars.
Because light has a finite speed, looking across large distances in the
Universe is equivalent to looking back in time. Therefore, the
observation of active galactic nuclei at large distances and their
scarcity in the nearby Universe suggests that they were much more common
in the early Universe, implying that active galactic nuclei could be early stages of galactic evolution.
This leads to the question about what would be the local (modern-day)
counterparts of AGNs found at large redshifts. It has been proposed that
NLSy1s could be the small redshift counterparts of quasars found at
large redshifts (z>4). The two have many similar properties, for
example: high metallicities or similar pattern of emission lines (strong Fe [II], weak O [III]).
Some observations suggest that AGN emission from the nucleus is not
spherically symmetric and that the nucleus often shows axial symmetry,
with radiation escaping in a conical region. Based on these
observations, models have been devised to explain the different classes
of AGNs as due to their different orientations with respect to the
observational line of sight. Such models are called unified models.
Unified models explain the difference between Type I and Type II
galaxies as being the result of Type II galaxies being surrounded by
obscuring toruses which prevent telescopes from seeing the broad line
region. Quasars and blazars can be fit quite easily in this model.
The main problem of such an unification scheme is trying to explain why
some AGN are radio loud while others are radio quiet. It has been
suggested that these differences may be due to differences in the spin
of the central black hole.
Centaurus A or NGC 5128,
apparently the brightest Seyfert galaxy as seen from Earth; a giant
elliptical galaxy and also classified as a radio galaxy notable for its relativistic jet spanning more than a million light years in length.
Cygnus A, the first identified radio galaxy and the brightest radio source in the sky as seen in frequencies above 1 GHz
Messier 51a (NGC 5194), the Whirlpool Galaxy, one of the best known galaxies in the sky
The rapid neutron-capture process, or so-called r-process, is a set of nuclear reactions that in nuclear astrophysics is responsible for the creation of approximately half of the atomic nucleiheavier than iron; the "heavy elements", with the other half produced by the p-process and s-process. The r-process usually synthesizes the most neutron-rich stable isotopes of each heavy element. The r-process
can typically synthesize the heaviest four isotopes of every heavy
element, and the two heaviest isotopes, which are referred to as r-only nuclei, can only be created via the r-process. Abundance peaks for the r-process occur near mass numbersA = 82 (elements Se, Br, and Kr), A = 130 (elements Te, I, and Xe) and A = 196 (elements Os, Ir, and Pt).
The r-process entails a succession of rapidneutron captures (hence the name) by one or more heavy seed nuclei, typically beginning with nuclei in the abundance peak centered on 56Fe. The captures must be rapid in the sense that the nuclei must not have time to undergo radioactive decay (typically via β− decay) before another neutron arrives to be captured. This sequence can continue up to the limit of stability of the increasingly neutron-rich nuclei (the neutron drip line) to physically retain neutrons as governed by the short range nuclear force. The r-process therefore must occur in locations where there exist a high density of free neutrons. Early studies theorized that 1024 free neutrons per cm3
would be required, for temperatures about 1GK, in order to match the
waiting points, at which no more neutrons can be captured, with the
atomic numbers of the abundance peaks for r-process nuclei. This amounts to almost a gram of free neutrons in every cubic centimeter, an astonishing number requiring extreme locations. Traditionally this suggested the material ejected from the reexpanded core of a core-collapse supernova, as part of supernova nucleosynthesis, or decompression of neutron-star matter thrown off by a binary neutron star merger. The relative contributions of these sources to the astrophysical abundance of r-process elements is a matter of ongoing research.
A limited r-process-like series of neutron captures occurs to a minor extent in thermonuclear weapon explosions. These led to the discovery of the elements einsteinium (element 99) and fermium (element 100) in nuclear weapon fallout.
The r-process contrasts with the s-process, the other predominant mechanism for the production of heavy elements, which is nucleosynthesis by means of slow captures of neutrons. The s-process primarily occurs within ordinary stars, particularly AGB stars, where the neutron flux is sufficient to cause neutron captures to recur every 10–100 years, much too slow for the r-process, which requires 100 captures per second. The s-process is secondary,
meaning that it requires pre-existing heavy isotopes as seed nuclei to
be converted into other heavy nuclei by a slow sequence of captures of
free neutrons. The r-process scenarios create their own seed
nuclei, so they might proceed in massive stars that contain no heavy
seed nuclei. Taken together, the r- and s-processes account for almost the entire abundance of chemical elements heavier than iron. The historical challenge has been to locate physical settings appropriate for their time scales.
History
Following pioneering research into the Big Bang and the formation of helium in stars, an unknown process responsible for producing heavier elements found on Earth from hydrogen and helium was suspected to exist. One early attempt at explanation came from Chandrasekhar and Louis R. Henrich who postulated that elements were produced at temperatures between 6×109 and 8×109K. Their theory accounted for elements up to chlorine, though there was no explanation for elements of atomic weight heavier than 40 amu at non-negligible abundances.
This became the foundation of a study by Fred Hoyle,
who hypothesized that conditions in the core of collapsing stars would
enable nucleosynthesis of the remainder of the elements via rapid
capture of densely packed free neutrons. However, there remained
unanswered questions about equilibrium in stars that was required to
balance beta-decays and precisely account for abundances of elements that would be formed in such conditions.
The need for a physical setting providing rapid neutron capture,
which was known to almost certainly have a role in element formation,
was also seen in a table of abundances of isotopes of heavy elements by Hans Suess and Harold Urey in 1956. Their abundance table revealed larger than average abundances of natural isotopes containing magic numbers of neutrons as well as abundance peaks about 10 amu lighter than stable nuclei
containing magic numbers of neutrons which were also in abundance,
suggesting that radioactive neutron-rich nuclei having the magic neutron
numbers but roughly ten fewer protons were formed. These observations
also implied that rapid neutron capture occurred faster than beta decay, and the resulting abundance peaks were caused by so-called waiting points at magic numbers. This process, rapid neutron capture by neutron-rich isotopes, became known as the r-process, whereas the s-process
was named for its characteristic slow neutron capture. A table
apportioning the heavy isotopes phenomenologically between s-process and r-process isotopes was published in 1957 in the B2FH review paper, which named the r-process and outlined the physics that guides it. Alastair G. W. Cameron also published a smaller study about the r-process in the same year.
The stationary r-process as described by the B2FH paper was first demonstrated in a time-dependent calculation at Caltech by Phillip A. Seeger, William A. Fowler and Donald D. Clayton, who found that no single temporal snapshot matched the solar r-process abundances, but, that when superposed, did achieve a successful characterization of the r-process abundance distribution. Shorter-time distributions emphasize abundances at atomic weights less than A = 140, whereas longer-time distributions emphasized those at atomic weights greater than A = 140. Subsequent treatments of the r-process reinforced those temporal features. Seeger et al. were also able to construct more quantitative apportionment between s-process and r-process of the abundance table of heavy isotopes, thereby establishing a more reliable abundance curve for the r-process isotopes than B2FH had been able to define. Today, the r-process abundances are determined using their technique of subtracting the more reliable s-process isotopic abundances from the total isotopic abundances and attributing the remainder to r-process nucleosynthesis. That r-process
abundance curve (vs. atomic weight) has provided for many decades the
target for theoretical computations of abundances synthesized by the
physical r-process.
The creation of free neutrons by electron capture during the
rapid collapse to high density of a supernova core along with quick
assembly of some neutron-rich seed nuclei makes the r-process a primary nucleosynthesis process, meaning a process that can occur even in a star initially of pure H and He, in contrast to the B2FH designation as a secondary process
building on preexisting iron. Primary stellar nucleosynthesis begins
earlier in the galaxy than does secondary nucleosynthesis. Alternatively
the high density of neutrons within neutron stars would be available
for rapid assembly into r-process nuclei if a collision were to
eject portions of a neutron star, which then rapidly expands freed from
confinement. That sequence could also begin earlier in galactic time
than would s-process nucleosynthesis; so each scenario fits the earlier growth of r-process abundances in the galaxy. Each of these scenarios is the subject of active theoretical research.
Observational evidence of the early r-process enrichment of
interstellar gas and of subsequent newly formed of stars, as applied to
the abundance evolution of the galaxy of stars, was first laid out by
James W. Truran in 1981.
He and subsequent astronomers showed that the pattern of heavy-element
abundances in the earliest metal-poor stars matched that of the shape of
the solar r-process curve, as if the s-process component were missing. This was consistent with the hypothesis that the s-process had not yet begun to enrich interstellar gas when these young stars missing the s-process abundances were born from that gas, for it requires about 100 million years of galactic history for the s-process to get started whereas the r-process can begin after two million years. These s-process–poor, r-process–rich stellar compositions must have been born earlier than any s-process, showing that the r-process
emerges from quickly evolving massive stars that become supernovae and
leave neutron-star remnants that can merge with another neutron star.
The primary nature of the early r-process thereby derives from observed abundance spectra in old stars that had been born early, when the galactic metallicity was still small, but that nonetheless contain their complement of r-process nuclei.
Periodic table
showing the cosmogenic origin of each element. The elements heavier
than iron with origins in supernovae are typically those produced by the
r-process, which is powered by supernovae neutron bursts
Either interpretation, though generally supported by supernova
experts, has yet to achieve a totally satisfactory calculation of r-process
abundances because the overall problem is numerically formidable, but
existing results are supportive. In 2017, new data about the r-process was discovered when the LIGO and Virgo gravitational-wave observatories discovered a merger of two neutron stars ejecting r-process matter.
Noteworthy is that the r-process is responsible for our
natural cohort of radioactive elements, such as uranium and thorium, as
well as the most neutron-rich isotopes of each heavy element.
Immediately after the severe compression of electrons in a Type II supernova, beta-minus decay is blocked. This is because the high electron density fills all available free electron states up to a Fermi energy which is greater than the energy of nuclear beta decay. However, nuclear capture of those free electrons still occurs, and causes increasing neutronization of matter. This results in an extremely high density of free neutrons which cannot decay, on the order of 1024 neutrons per cm3), and high temperatures. As this re-expands and cools, neutron capture by still-existing heavy nuclei occurs much faster than beta-minus decay. As a consequence, the r-process runs up along the neutron drip line and highly-unstable neutron-rich nuclei are created.
Three processes which affect the climbing of the neutron drip line are a notable decrease in the neutron-capture cross section in nuclei with closed neutron shells, the inhibiting process of photodisintegration, and the degree of nuclear stability in the heavy-isotope region. Neutron captures in r-process nucleosynthesis leads to the formation of neutron-rich, weakly bound nuclei with neutron separation energies as low as 2 MeV. At this stage, closed neutron shells at N
= 50, 82, and 126 are reached, and neutron capture is temporarily
paused. These so-called waiting points are characterized by increased
binding energy relative to heavier isotopes, leading to low neutron
capture cross sections and a buildup of semi-magic nuclei that are more
stable toward beta decay.
In addition, nuclei beyond the shell closures are susceptible to
quicker beta decay owing to their proximity to the drip line; for these
nuclei, beta decay occurs before further neutron capture. Waiting point nuclei are then allowed to beta decay toward stability before further neutron capture can occur, resulting in a slowdown or freeze-out of the reaction.
Decreasing nuclear stability terminates the r-process when its heaviest nuclei become unstable to spontaneous fission, when the total number of nucleons approaches 270. The fission barrier may be low enough before 270 such that neutron capture might induce fission instead of continuing up the neutron drip line. After the neutron flux decreases, these highly unstable radioactive nuclei undergo a rapid succession of beta decays until they reach more stable, neutron-rich nuclei. While the s-process creates an abundance of stable nuclei having closed neutron shells, The r-process, in neutron-rich predecessor nuclei, creates an abundance of radioactive nuclei about 10 amu below the s-process peaks after their decay back to stability.
The r-process also occurs in thermonuclear weapons, and
was responsible for the initial discovery of neutron-rich almost stable
isotopes of actinides like plutonium-244 and the new elements einsteinium and fermium
(atomic numbers 99 and 100) in the 1950s. It has been suggested that
multiple nuclear explosions would make it possible to reach the island of stability,
as the affected nuclides (starting with uranium-238 as seed nuclei)
would not have time to beta decay all the way to the quickly spontaneously fissioning nuclides at the line of beta stability before absorbing more neutrons in the next explosion, thus providing a chance to reach neutron-rich superheavy nuclides like copernicium-291 and -293 which should have half-lives of centuries or millennia.
Astrophysical sites
The most probable candidate site for the r-process has long been suggested to be core-collapse supernovae (spectral types Ib, Ic and II), which may provide the necessary physical conditions for the r-process. However, the very low abundance of r-process nuclei
in the interstellar gas limits the amount each can have ejected. It
requires either that only a small fraction of supernovae eject r-process nuclei to the interstellar medium, or that each supernova ejects only a very small amount of r-process
material. The ejected material must be relatively neutron-rich, a
condition which has been difficult to achieve in models, so that astrophysicists remain uneasy about their adequacy for successful r-process yields.
In 2017, entirely new astronomical data about the r-process was discovered in data about the merger of two neutron stars. Using the gravitational wave data captured in GW170817 to identify the location of the merger, several teams observed and studied optical data of the merger, finding spectroscopic evidence of r-process
material thrown off by the merging neutron stars. The bulk of this
material seems to consist of two types: hot blue masses of highly
radioactive r-process matter of lower-mass-range heavy nuclei (A < 140 such as Strontium) and cooler red masses of higher mass-number r-process nuclei (A > 140) rich in actinides (such as Uranium, Thorium, and Californium).
When released from the huge internal pressure of the neutron star,
these ejecta expand and form seed heavy nuclei that rapidly capture free
neutrons, and radiate detected optical light for about a week. Such
duration of luminosity would not be possible without heating by internal
radioactive decay, which is provided by r-process nuclei near their waiting points. Two distinct mass regions (A < 140 and A > 140) for the r-process yields have been known since the first time dependent calculations of the r-process.
Because of these spectroscopic features it has been argued that such
nucleosynthesis in the Milky Way has been primarily ejecta from
neutron-star mergers rather than from supernovae.
These results offer a new possibility for clarifying six decades of uncertainty over the site of origin of r-process nuclei. Confirming relevance to the r-process is that it is radiogenic power from radioactive decay of r-process nuclei that maintains the visibility of these spun off r-process fragments. Otherwise they would dim quickly. Such alternative sites were first seriously proposed in 1974 as decompressing neutron star matter. It was proposed such matter is ejected from neutron stars merging with black holes in compact binaries. In 1989 (and 1999) this scenario was extended to binary neutron star mergers (a binary star system of two neutron stars that collide). After preliminary identification of these sites, the scenario was confirmed in GW170817. Current astrophysical models suggest that a single neutron star merger event may have generated between 3 and 13 Earth masses of gold.
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