Introduction

Misc. Galactic Variables

Introduction

These stars defy categorization with any of the large groups of variable stars. Some of the peculiar variable stars are close to fitting into some known group, but others like those discussed here are, as far as is known, completely unique.


R Corona Borealis

R Corona Borealis variables form a numerically small but distinctive group of stars. Hey cannot be considered either pulsating variables, as eruptive variables or as eclipsing binaries, so that they form a separate class. Their normal state is at the maximum, but at irregular and generally widely space intervals; the brightness rapidly decreases by 4, 6, or 8 magnitudes. The reversion to the normal state sometimes occurs in a few weeks but can also take several years.

The prototype star, R Corona Borealis, was discovered in 1795. It is normally of magnitudes 6 with small fluctuations, but at its minima it can fall to below magnitude 14. The decrease is paid (1 or 2 months) but the rise can take a long time. The most remarkable period lasted from 1864 to 1875; during these 11 years, R Corona Borealis did not return to its maximum but underwent several large fluctuations. There are few, 2 or 3 dozen, of this type of star, several being outside the galaxy (W Men, in the Large Megallanic Cloud).

These stars are supergiants with G spectra and absolute magnitude around -6. Their spectra, which have been subjected to considerable study, are normally rich in carbon and poor in hydrogen. At a minimum, there are several metallic lines in emission (neutral sodium, ionized calcium, titanium, and iron) but there are absorption lines of hydrogen. In addition, and this is an important, the radial velocity varies very little between minima and maxima, which rules out the possibility that matter is being ejected.

A semi-periodic pulsation has shown to exist in the some of these stars, but only of small amplitude (0.1-0.3 magnitude). These stars generally have a small color index; for R CrB, the B-V index is only +0.60, slightly lower than the Sun (+0.65). Observations in the infrared show that they are much brighter in that part of the spectrum than in the visible region. For example RY Sgr, which has a visual magnitude of 6.5 at maximum, reaches a magnitude of -0.4 in the far infrared where it is brighter than Vega.

The interpretation of all these observations runs into difficulties. The stars are usually rich in carbon and it is thought that this element plays the same role in stars that water vapor in the Earth’s atmosphere, condensing to form clouds that obscure the sky unless they disappear in the form of rain.

The carbon circulating in the star’s atmosphere condenses in the upper regions to form fine grains of graphite. During their movement towards the surface these clouds of carbon become denser and form a think envelope that obscures the star and cuts out a large part of its radiation. However as the clouds get even nearer to the surface their temperature rises and the carbon again becomes gaseous by sublimation. The atmosphere is clear once again and everything is in order waiting for the next cycle to begin. One question regarding this theory is how the carbon is ejected into the upper atmosphere.


Unique Variable Stars

V389 Cyg

There are some variable that are catalogued as “Unique Variables” this is an extremely small number of stars and also some of these are very close to being able to be catalogued as a known type. There remain a few that do not fit any known type of variable.

V389 Cyg is a star that oscillates between magnitude 5.5 and 5.69 with two periods: the first is 1.12912 days with an amplitude of 0.19 magnitude. The second period is 1.19328 days with an amplitude of 0.05. These periods do not superpose as they would if it were a normal pulsating star. For a time which can be anywhere from 1 to 3 weeks, the star varies with period 1, and after that it varies with period 2. Sometimes in between the two periods it stops varying at all.

Adding more complexity, the radial velocity varies with yet a third period of 3.31322 days. Note this is not a multiple of the other two periods. The maximum radial velocity sometimes coincides with the maximum brightness. This is what happens with the Cepheids but V389 Cyg is a B8 star which means that it cannot be a Cepheid.

AE Aqr

AE Aqr is a binary consisting of a red supergiant with an emission spectrum of K5 IVe and a red subdwarf sdB. It is a well-studied spectroscopic binary of period 0.4116 day. Photoelectric observations have shown some odd photometric variations: flares are observed similar to those of UV Ceti stars but with superposed rapid irregular variations, as in dwarf nova and ex-nova. In addition the flares are larger when the star is brighter.

The red component is undoubtedly responsible for the flares and the B star for the flickering, but the fact that there is a correlation between the flares and the flickering shows that the two activities are closely connected. No one currently knows how.

V384 Cyg

V384 Cyg varied between magnitude 10.6 and 17 from 1907 to 1927 and its light curve was unhesitatingly classed among the R Carona Borealis star. Since 1927, the amplitude has been much smaller and the variations have become semi-periodic with a cycle of about 200days. Its spectrum is very odd: at maximum brightness, only emission lines can be seen, with neutral helium, ionized oxygen and ionized carbon, but no hydrogen lines. At the minimum, there is still oxygen and now hydrogen in emission also: carbon on the other hand, has disappeared. The star is also surrounded by faint nebulosity. According to Gerbig, it is probable that the luminous variations are due to the ejection of an opaque envelope. However, it is still impossible to understand why the hydrogen and carbon disappear and reappear alternatively.

V605 Aql

This star is normally very faint, with a photographic magnitude of 20, but it had a maximum accompanied by large fluctuations that took it to magnitude 11 between 1918 and 1924. Some observations have purported to show that the light curve is similar to that of the slow supernova in NGC 1058. This star is not a supernova because in 1921, at the time of its maximum t was found that it had a carbonized spectrum similar to that of R Coronae Borealis! In addition it was surrounded by a small planetary nebula, A58. It is therefore rather to be classed among the novoids, but further observations are needed.

V1016 Aql

Between 1920 and 1963, the star V1016 Cyg behaved like a long period variable, oscillating between the magnitudes of 15.0 and 17.5 with a period of 450 days. From 1964, its brightness began to increase and in 1972 it reached magnitude of 10.3. Since then this has maintained, except for small fluctuation. From 1965, its spectrum showed extremely strong emission lines and in 1974 a small planetary nebula was seen to appear and increase rapidly in size We may add that, although the Mira-type variations disappeared in the visible region, they persisted in the infrared and the period was unchanged at 450 days.

V1329 Aql

V1239 Cyg is a star with a M5 spectrum. Between 1891 and 1964, it varied irregularly between magnitudes 14 and 16; since 1965, the brightness has increased to 11.5 and superposed on the M spectrum there is now a planetary nebular spectrum with very strong emission. In addition, V12329 Cyg displays eclipses with a period of 959 days and amplitude of 2 magnitudes.

FG Sagittae

A white dwarf in the making

Variable stars change more than most, and more noticeably, too. Stellar evolution, the physical changes that stars undergo from birth to death, is far from completely understood. Because many of these processes take millions or billions of years, we have to infer what changes happen inside stars over time from bits and pieces of evidence. But a few stars are doing interesting things on much shorter timescales, and one of them is our Variable Star of the Season -- the highly-evolved star FG Sge, a white dwarf in the making.

FG Sge was first noted as a variable in 1944 by Cuno Hoffmeister, but the star received scant observer attention (at least in printed papers), until 1960, when G. Richter published a study of the long-term photometric and spectroscopic behavior of FG Sge (or "377.1943 Sge"). The brightness of FG Sge increased by over four magnitudes since the start of the century, and he found evidence of smaller, short-term variations as well. Richter saw signatures of a nova in objective prism spectra, and he likened the star to the Z Andromedae-type symbiotic star AG Pegasi. Since then, many papers have detailed the remarkable evolution of FG Sge from a faint, hot, blue post-asymptotic giant branch star and planetary nebula in the making to a much cooler and brighter yellow supergiant. Even more exciting for variable star observers, following FG Sge's four-magnitude brightening and several decades of relative constancy, the star now appears to exhibit the dramatic and seemingly random fluctuations and fadings of the R Coronae Borealis class of variable stars. FG Sge is clearly a star undergoing extreme changes, and we're fortunate to be treated to its amazing show.

Stellar Evolution: from red giants to white dwarfs (and back again?)

When stars like the Sun near the ends of their lives, they begin to evolve more quickly and more dramatically than the slow, steady evolution that occurs during the main sequence. When the nuclear burning core of a Sun-like star runs out of the hydrogen that powers main sequence evolution, the hot sphere of helium left behind shrinks and heats, and a layer of unprocessed hydrogen around this core can then ignite -- a red giant is born. From here, the evolution of low- and intermediate-mass stars becomes very complicated, involving several episodes of
• the rapid nuclear burning of helium in the core to carbon (in a "helium flash"),

• mass loss,

• the dredge-up of core material to the surface by convection

• very strong pulsations (as in the Mira stars), and, at the very end

• the expulsion of the outer layers of the star, leaving behind a fading planetary nebula with an inert white dwarf star at its heart.

The pace at which these changes occur becomes increasingly rapid; main sequence evolution may take billions of years, red giant evolution may take many millions of years, and AGB evolution may take a few million years. But the time from the AGB to the planetary nebula stage is measured in thousands of years, and some very large changes may occur of a span of decades, observable in a single human lifetime. Although there are billions of stars in our Galaxy, the end stages of a star's life are so short that we're lucky if we get to see just a few stars exhibiting these changes at any given time. We're pretty lucky indeed to see a spectacular show like that of FG Sge!

A star in transition

In their excellent 2006 review paper, C.S. Jeffery and D. Schönberner assessed what we know (and think we know) about this amazing star, comparing the many years and many layers of observation and analysis to an "archaeological dig" -- using the historical record to understand the past and predict the future. And FG Sge has had quite a history. At the time of its discovery, FG Sge (or "377.1943 Sge") was what is known as a post-asymptotic giant branch star or "post-AGB star". It had evolved through most of the thermal pulses that occur on the asymptotic giant branch and had begin to blow off its outer layers as a planetary nebula (Hen 1-5, discovered by Karl Henize in 1961). The star probably began to brighten in the late 1800's as a result of a late thermal pulse, the nuclear ignition of a thin helium shell surrounding the inert carbon-oxygen core of the star. As with thermal pulses on the AGB, the result of this was that the post-AGB star brightened considerably as its outer layers expanded due to the enormous amount of energy released by the shell. Somewhat counterintuitively, like all giant stars, FG Sge became cooler as its luminosity was distributed over a much, much larger surface area. As a result, FG Sge became much brighter but also much redder over several decades. Prior to the dramatic fading experienced in the early 1990's, it was classified as a yellow supergiant rather than a faint, blue proto-white dwarf. That's quite a change in less than a century!

Another change that stars undergo late in life is a change in chemical composition. On one hand, stars begin to lose mass from their outer layers, and much of what is lost at first has the relatively unchanged chemical composition the star had at its birth -- mostly hydrogen and helium. A larger and larger fraction of the remaining mass of the star is then composed of material that has undergone nuclear fusion in the core. On the other hand, the physical process of convection is very strong in cool, giant stars, and during the end stages of a star's life, convection can even dredge up material from all the way down in the core of the star. It was once thought FG Sge might be in the process of dredging up some of this processed nuclear material. The extreme chemical changes suggested in earlier studies of FG Sge are probably not real, but it is clear that the star is changing.

AAVSO light curve of FG Sge. Black points are visual observations; green are V, blue are B, and red are Rc-band observations. FG Sge declined sharply in brightness in 1992, much like the R CrB stars; it is believed the star became dust enshrouded at this time. The R CrB-like fadings and rebrightenings have continued to the present.

Like many AGB stars, FG Sge already shows enhanced spectral signatures of many metallic elements known as s-process or "slow process" elements, those that can be generated by slow neutron capture. (The AGB stars are the originators of the s-process elements we see on Earth today.) Beginning in the 1960's and continuing until relatively recently, there were a number of spectroscopic studies of FG Sge that suggested the star might be in the process of mixing nuclear-processed core material to the surface as we're watching. Since then, very careful analyses of stellar atmosphere models, post-AGB star evolution models, and much photometric and spectroscopic data now shows that such a recent enhancement of s-process elements has not occurred. However, things are changing. The hydrogen abundance at the surface is decreasing, most likely because of mass loss -- it is after all a planetary nebula in the making! The helium abundance at the surface may also be increasing, and this may be a sign that a dredge-up event is underway, and has just reached the fringes of the core. Confirmation of that may require several decades (or centuries) of observations, including those by variable star observers like you.

Although it's not clear that the overall chemistry is changing right now, it certainly did change in the not-too-distant past. One important change caused by a past dredge-up event was an enhancement of the element carbon relative to oxygen. When the carbon abundance exceeds that of oxygen (C/O > 1), the star becomes a carbon star, much redder than other AGB stars where the oxygen abundance is less than or comparable to that of carbon. Because FG Sge has also been cooling, the increased carbon abundance allows the formation of carbon (C2) molecules, which are an important component of dust around AGB stars. Indeed, as many of you discovered for yourselves around 1990, FG Sge suddenly became dust enshrouded, leading to a dramatic decrease in brightness similar to those observed in the R Coronae Borealis stars. These irregular fadings have continued, and FG Sge remains faint right now. It's an open question of how long this behavior will continue, or whether observers will be in for even more changes in the coming years and decades. But it's likely that FG Sge will continue to show dramatic changes in behavior until its nuclear-burning helium shell finally runs out of gas, a process that may take several centuries. As with many AGB stars, your observations -- and those of observers who will come after -- may well be used by future astronomers and astrophysicists to understand this curious star. The cosmic "archaeology" that Jeffery and Schönberner undertook in their paper will continue well into the future!

The young planetary nebulae Hen 2-47 (left) and Hen-1357 (right), as imaged by the Hubble Space Telescope. FG Sge is already in the earliest stages of forming such a nebula itself. (Credit -- He 2-47: NASA, ESA, and The Hubble Heritage Team (STScI/AURA); Hen-1357: Matt Bobrowsky (OSC) and NASA)

Rare, but not alone

There are a few other stars thought to be nearing the end of their lives, too, including V605 Aql and the famous Sakurai's Object -- V4334 Sgr, which was itself a topic of our Variable Star of the Month series. Both of these objects seem to show what FG Sge will be like many centuries or millenia hence. V605 Aql seems to once again be an extremely hot star, and is settling back onto the post-AGB or white dwarf cooling track after one last thermal pulse. And the spectacular Sakurai's Object might just be a star that experienced this last very late thermal pulse (VLTP) while we watched. Indeed, it was considered a slow nova when it was first discovered, but it is now clear it is a single star rather than a typical binary nova progenitor. Both stars have faded back to near invisibility in the optical since their discovery. Although rare, FG Sge does have at least a little company in the Milky Way!

A notable counterexample is V838 Mon, another supposed nova from 2002. Although this famous object with its stunning light echo remains somewhat mysterious, it is most likely not a single evolved star but an even rarer event -- a stellar merger. The rapid increase in brightness was likely caused by the merger itself, while the surrounding dust shells may have come from mass loss by one of the pre-merger stars if it was highly evolved. Such cases of stellar mergers are even rarer than catching a star at the end of its life -- an important exception to the rule when you see hoofprints, think horses, not zebras!

Observing FG Sagittae

FG Sge is a challenging object to observe. It's very faint now, below 15th magnitude, and lies within a crowded Milky Way field. It is also surrounded by nebulosity -- the planetary nebula it ejected prior to its current activity -- and while it should be easy to pick out in the field, it won't be so easy to photometer! When you obtain charts via VSP, we recommend you create them with DSS images rather than plotting points; and make sure you report the brightness of the central star, rather than the entire nebula! The integrated light of the nebula is much brighter than FG Sge itself right now. Both visual and CCD observations are needed, including fainter-thans. Many of the most recent observations submitted are fainter-thans, and these will be useful for long-term studies; knowing that the star was fainter than 15 is a good constraint on its behavior, even if we don't know what the actual magnitude was.

If you're a CCD observer and can do multicolor photometry, try taking images in all of the photometric filters you have available. What color is FG Sge? What color is the integrated light of the nebulosity? The nebula is approximately 30 arcseconds across, so it will be resolved in most telescopes. Try to make a multicolor image if you have the software -- photometry is important, but the field will make a lovely astrophoto as well. If you get a good image, let us know, we'd love to see it!

FG Sge is observable all throughout the northern hemisphere, and from most populated areas in the south. It lies within the Milky Way (at a Galactic latitude of -7°), which itself makes for a beautiful target at this time of year. It's near the eastern edge of the constellation Sagitta, a small but prominent arrow-shaped constellation sandwiched between Vulpecula (to the north) and Aquila (to the south). Look for the Galactic globular cluster M71 not too far to the west.

FG Sge — a puzzle for present and future

Our Variable Star of the Season, FG Sge, will be doing interesting things for many years -- if not centuries -- to come, and is sure to provide an interesting show for variable star observers for a long, long time. Stars like FG Sge are beautiful, but they teach a good lesson as well. Even things as giant and stately as the stars in the sky come to an end, and we're fortunate to be witnessing one such event in our lifetimes. Although the final death throes of FG Sge are less dramatic than the supernovae that come to mind when we think of the passing of stars, its end will have no less spectacular consequences. FG Sge itself will become an another anonymous white dwarf among billions in our galaxy, but its ashes will be swept up into new generations of stars, and maybe even new generations of star gazers, too. We hope you get many chances this season to enjoy our lovely Milky Way, and perhaps catch a glimpse of FG Sge's swan song.


Rotating variable stars

Another class of variable stars owes their change in brightness to their irregular shape and/or to their non-uniform surface. In some cases, the star's shape may not be a perfect sphere, but rather an ellipsoid. Alternatively, the brightness across the surface of the star itself may vary because of temporarily brighter and darker patches on the star's surface; similar to sunspots on the Sun (in this case they are called starspots).

These stars are classed as rotating variables. The observed change in their apparent brightness is due to the rotation of the star. This means that brighter and dimmer areas come into view as the star turns (in the case of starspots) or different amounts of the star's surface area (in the case of stars that are not perfectly spherical).

Rotating variable star type: Ellipsoidal

These rotating variable stars fall into the category of stars that are not perfectly spherical in shape, but are an ellipsoidal. Ellipsoidal variables stars are often in a binary system where the two stars that are orbiting around one another are very close together. These binary stars are so close together that their mutual gravitational attraction distorts their shape.

The components have an ovoid shape but the plane of revolution makes a large angle with the direction of our line of site. There are therefore no eclipses, but merely a variation in the apparent stellar area. This produces a slight variation in the brightness (a few hundredths of a magnitude) detectible in bright stars. A few of these stars are found in the table.

These are all stars of B or A spectra concentrated in the galactic arms (Orion and Perseus), and they appear in OB associations or young galactic clusters. They therefore form a typical Population I.

Some Ellipsoidal Variables
ObjectPeriod
(days)
Eta-5 5 Ori 3.70
Psi-2 Ori 2.526
IW Per 0.917
IX Per 1.326
Omicron Per 4.419
Beta Per 1.527

The New Variable Stars

Introduction


X-ray Binaries

X-ray binaries encompass a large class of objects involving neutron stars or black-hole candidates. Neutron stars are created in supernovae, and they are generally more massive than the Sun, with maximum masses around 2.5 solar masses—the precise mass depends on the physics of the interior, which is highly uncertain. Above this upper limit, the compact object under the theory of general relativity must be a black-hole because the pressure at the center of such a massive star cannot counteract the force of gravity.

The free-fall velocity at the surface of a neutron star is about half of the speed of light. The gravitational potential energy released when matter free-falls to the surface of a neutron star is therefore about 15% of the rest mass energy, which is much more than the 0.8% that is released in converting hydrogen into iron. The gravitational potentials of neutron stars and black-hole candidates provide the most effective mechanism of releasing energy from matter. But while the maximum possible temperature that can be achieved through this mechanism is a substantial fraction of the rest-mass energy of the matter, which for hydrogen places the maximum temperature at several hundred MeV, the characteristic temperatures seen in the radiation is in the tens to hundreds of keV range, so that the energy appears as x-rays and gamma-rays.

For neutron stars, about half of the gravitational potential energy is released as matter flows through the accretion disk, and the remainder is released when the matter falls from the accretion disk into the neutron star's atmosphere. If the black-hole candidates are in fact black holes, the gravitational potential energy is released only through the accretion disk, because black holes do not have a surface that matter can strike, but an event horizon towards which matter, from the standpoint of an observer far from the system, slowly settles.

In an accretion disk, the matter orbits the compact object many times as it slowly drifts to the compact objects. The viscosity in the accretion disk slowly converts gravitational energy into thermal energy. This slow heating of the matter is counteracted by the release of electromagnetic radiation; the temperature of the disk adjusts itself to balance the cooling against the heating. The accretion disk pages describe how thepower released by an accretion disk is related to its minimum temperature. For a minimum temperature in the 1 keV range, the mass flow rate onto a solar-mass neutron star must be above 1038 ergs s-1, which makes these systems 104time more luminous than the Sun.

For a black hole, the inner edge of the accretion disk does not extend to the event horizon. Black holes have a peculiar property: circular orbits around black holes only exits outside a certain critical radius beyond the event horizon. This critical radius is called the point of last stable orbit. For an object to orbit in a circle at this radius, it must travel at the speed of light. Inside this point, every orbit that starts parallel to the event horizon bends and meets the event horizon. This means that as matter in an accretion disk flows beyond the last stable orbit, it fall towards the event horizon, reaching the event horizon before it can complete a full orbit. Once on this path, the viscosity that had converted gravitational potential energy into radiation ceases to work, so that the gravitational potential energy remains with the matter as it falls to the event horizon.

The accretion of material into the atmosphere of a neutron star produces at least as much energy as is released in the accretion disk. If the neutron star has a weak magnetic field, the accretion disk extends to the neutron star's atmosphere, where the problem becomes very difficult from a theoretical standpoint. The interaction is between supersonic (half the speed of light) centrifugally-supported fluid and a pressure-supported subsonic atmosphere. What can be said from the temperatures seen in these systems is that the temperature characterizing the radiation is much lower than the temperatures encountered in a shock.

If the magnetic field on the neutron star is strong enough, it disrupts the inner-edge of the accretion disk before it reaches the star's atmosphere, and the material in the disk flows along magnetic field lines to the magnetic poles of the star, making the magnetic poles x-ray hot-spots on the neutron star. As the neutron star rotates, these hot spots change orientation or pass behind the star, making the x-ray luminity of such a star oscillate at the star's rotation period. These systems are called x-ray pulsars.

As stated above, most of the energy released from a neutron-star binary is gravitational potential energy. But nuclear fusion still plays a role in these systems. Normally the neutron star is accreting a mixture of hydrogen and helium onto its surface from its companion. This means that the atmosphere of an accreting neutron star is composed of hydrogen and helium under extreme pressures. This leads to the nuclear fusion of hydrogen and helium into heavier elements. In some systems, nuclear fusion burns hydrogen and helium as fast as it is added to the star, so that the energy of nuclear fusion is a small addition to the continuous release of gravitational potential energy, but in other systems, fusion cannot keep pace, and hydrogen and helium build up on the star until the star's atmosphere is unstable to a nuclear runaway. In a runaway, the lower part of the atmosphere of the star detonates, and this immediate release of nuclear energy exceeds the release of gravitational potential energy for a short time. This energy diffuses through the neutron star's upper atmosphere and is released as x-rays. These systems, which detonate regularly, are called x-ray bursters.

Neutron and More Massive Companion

Probes in the 1970’s first detected x-rays that were discovered to be coming from binary star systems. The first of such systems to be discovered was Cen X3 (V779 Cen) in 1971 by the Uhuru satellite. It exhibits an x-ray pulse of 4.842 seconds. The system also shows eclipses with a period of 2.08726 days, and during the eclipse, the pulsation stops; it is clearly the x-ray component that is eclipsed.

The mechanism of the x-ray emission is: a binary star such as Cen X3 is formed from a normal hot star of O or B spectrum and a neutron star. Material escapes from the normal star and flows towards its companion, ether because the former occupies the whole of its Roche lobe, or because of the stellar wind, is captured in part by the neutron star. When the gas falls on to the companion, part of the kinetic energy is transformed into heat energy; in other words, it forms very hot plasma on the neutron star and thus generates high-energy radiation, x-rays. The plasma is rapped by the very strong magnetic field of the hyper-dense star in the regions around the magnetic poles and this produces periodic pulses, similar in principle to those of pulsars.

Several systems of this type have been discovered, and details of four (plus one) of them are shown in the table, including their visible spectra and the x-rays power emitted in J.s. The x-ray power is always enormous, from 10^29 – 10^32 J/s. The total emission of the sun over all wavelengths is only 10^23 J/s, so the power generated by these binaries is several thousand times greater.

Some X-ray Binaries of Large Mass
ObjectOrbitalPulsation  DistanceLuminostiyMass
Designationperiod (days)period (sec)M(pg)Spectra(KPC)(J/s)(Sun = 1)
Cen X3 (V779 Cen)2.14.813.4O6.5 II8.04z10^3017 + 1.0
Vel X1 (GP Vel)9.02836.9B0.5 Ib1.4 1x10^2920.5 + 1.7
4U 1538-533 (QV Nor)3.7 529 14.5B0 I 656x10^29 20 + 2.0
SMC X13.90.713.3B0 Ib656x10^3116.2 + 1.0
Cyg X1 (V1357 Cyg) 5.6Irregular8.9 O9.5 Iab 2.52x10^30 30 + 6

The importance of binaries is that they allow the mass of the components to be determined. It can be seen that, in the last four cases in the table, the mass of the hyper-dense companion is between 1 and 2 solar masses. This fits in well with theory that puts the upper limit of these objects to 2.5 to 3.5 solar masses.

Black Hole and More Massive Companion

The last star in the table has been determined to be a black hole, after further years of study, thus giving yet another type of x-ray binary star. Cyg X1 (V1357 Cyg) is very different from the neutron x-ray binaries in the table. First, its companion does not pulsate but only shows a slight ultra-rapid flickering, and secondly because it has a mass between 5 and 6 solar masses (which is over the limit for neutron stars). This type of binary is observed to have very rapid pulsations in the x-ray region in contrast to the extremely regular pulsations in the neutron x-ray binary systems.

Neutron and Less Massive Companion

Another type of x-ray binary is one whose neutron star is of equal or greater mass than that of the other member of the system, which is a dwarf star. This is the type of x-ray binary that is represented by Sco X1 (V818 Sco), this is also the first x-ray source observed in the sky. It is a binary with a period of 0.79 day; it has no pulsations but does show a continuous rapid flickering and sometimes some actual flares, detectable not only in the x-ray region but also in the visible region. This is a relatively close star at 700 parsecs and its x-ray luminosity is 2X10^30 J/s.


Polars: Magnetic Cataclysmic Variables

Introduction

The exotic star AM Herculis is the namesake of the "AM Her stars" or "polars", a unique class of cataclysmic variables in which the magnetic field of the primary star (white dwarf) completely dominates the accretion flow of the system. With the discovery of AM Her comes the discovery of "polars" and a lesson learned that even familiar objects will reveal exciting discoveries if they are approached in the correct way. AM Her was discovered in 1923 by M. Wolf in Heidelberg, Germany during a routine search for variable stars. It was then listed in the General Catalogue of Variable Stars as an irregular variable with a range from 12th to 14th magnitude. The listing remained as such until 1976, when the true complexities of AM Her were finally revealed. Berg & Duthie of the University of Rochester (1977) initially suggested that AM Her could be the optical counterpart of the weak X-ray source 3U 1809+50 which was detected by Uhuru, the first Small Astronomy Satellite. They noted that the variable star lay just outside of the region of certainty where the weak X-ray source was believed to be. Subsequently, a better position for 3U 1809+50 was determined and the position of the X-ray source and the variable star were shown to be the same. In order to prove that these observations were coming from the exact same source, however, more evidence was needed.

In May 1975, Berg and Duthie made the first photoelectric observations of AM Herculis. They found that the light from AM Her "flickered incessantly". This rapid light variability had been seen in two other stars which were associated with X-ray sources, so the team was optimistic that AM Her would turn out to be the optical counterpart of 3U 1809+50.

By May 1976, word had spread throughout the astronomical community that AM Herculis was an important object to observe in as much detail as possible. The Chilean astronomer S. Tapia, at the University of Arizona, had access to a polarimeter and used it to observe this star. The results were startling. He discovered in August 1976 that AM Her is both linearly and circularly polarized (these concepts are discussed later in the article) at optical wavelengths (Tapia 1977a). The detection of variable circular polarization was surprising since it was only known to exist in 9 other stars, all magnetic white dwarfs. The circular polarization in AM Her signified the presence of an immense magnetic field. This confirmed the suspicion that AM Her was the optical counterpart to the X-ray source. Consequently, a new class of magnetic cataclysmic variables called "AM Her stars" or "polars" was born.

Magnetic Cataclysmic Variables

The discovery of AM Herculis introduced a new class of highly magnetic stars to the group of cataclysmic variables known at the time. A cataclysmic variable is a close binary system with a white dwarf primary and a red dwarf secondary. Due to the evolution of the system, the red main sequence star loses matter in the direction of the primary, forming an accretion disk around the white dwarf. A magnetic cataclysmic variable is distinguished by the presence of a magnetic field around the white dwarf star that radically affects the whole nature of accretion flow in the system. Thus the cataclysmic variables are divided into two groups, the non-magnetic group (dwarf novae, nova-like, recurrent novae; for further review of non-magnetic CVs visit VSOTM for SS Cyg, U Gem, Z Cam, or RS Oph), and the magnetic group (polars). Magnetic cataclysmic variables are divided further into two classes based on the strength of their magnetic fields:

1. Intermediate Polars (DQ Her stars) The intermediate polars or DQ Her stars (named after the prototype DQ Her) show magnetic field strengths around the white dwarf star on the order of 1-10 Mega Gauss. An accretion disk forms, but is disrupted close to the white dwarf (primary) star due to the magnetic field. The magnetosphere is not strong enough to synchronize the orbits of the rotating white dwarf with the orbital period of the system (as seen in AM Her stars).

2. Polars (AM Her stars) The polars or AM Her stars (named after the prototype AM Her) display magnetic field strengths on the order of 10-100 Mega Gauss. This magnetic field is so powerful that it prevents the formation of an accretion disk around the white dwarf and locks the two stars together so they always present the same face to each other. Thus the white dwarf star spins at the same rate as the two orbit each other - a synchronous rotation that is the defining characteristic of an AM Her star. (About 10% of AM Her stars are asynchronous, where the rotations of the white dwarf and the orbit are off by ~ 1%) (Hellier 2001)

AM Her stars are especially interesting to study because of their strong magnetic fields. In an AM Her system the magnetic field around the white dwarf primary is so strong that no accretion disk is able to form like it does in non-magnetic cataclysmic variables. Material from the secondary flows toward the primary until it reaches the point where the magnetic field dominates the system. At this point, the energy associated with the magnetic field lines is much greater than the energy of the bulk flow of material coming from the secondary star so the material is forced to follow the path set by the field lines. This is usually a dipolar pattern that is similar to the configuration obtained by scattering iron filings around a bar magnet. Thus, to follow a field line the accretion stream splits into two, one part heads for the "north" magnetic pole and the other for the "south" pole. The field lines converge as they approach the white dwarf, squeezing the streams of matter and funneling them onto tiny accretion spots near the poles, whose radii are only ~ 1/100th of the white dwarf star (Hellier 2001). Liller (1977) describes the funneling of matter onto the magnetic poles of the white dwarf as resembling a "superviolent tornado". The flow of material to the magnetic poles is also similar to the aurora phenomenon on Earth, where solar particles enter Earth's atmosphere at those regions.

The material in this funnel, or accretion column, is channeled by the magnetic field towards the white dwarf in virtual free-fall. The potential energy is converted to kinetic energy and the stream slams into the white dwarf at roughly the ~ 3000 km/s escape velocity. In the resulting accretion shock the kinetic energy is converted into X-rays and radiated away. Magnetic cataclysmic variables emit most of their energy as X-rays and extreme-ultraviolet photons (Hellier 2001).

It is found that in an effort to obtain the lowest-energy configuration for this system, the magnetic field of the white dwarf often tilts over, so that one magnetic pole "points" towards the direction from which the stream comes. As a result, material flows preferentially onto that pole; material can still flow to the other pole, but only by going the long way around and only a fraction of the material succeeds in doing this. Eclipses in AM Her systems provide a graphic illustration of this stream geometry. Lightcurves reveal that the tiny and thus rapidly eclipsed, accretion spot at the magnetic pole emits roughly half of the total light; the other half comes from the extended stream, which enters and leaves eclipse more gradually (Hellier 2001).

A Look at the Light Curve

The light curve of AM Her appears to have the temperament of a superviolent tornado itself. There is apparently more than one source of radiation wreaking havoc on our star. The variations in AM Her may be thought to belong to two groups, the long-term changes and the short-term changes. The long-term changes are characterized by the existence of two different states, one the "active" or "on" state, in which the luminosity fluctuates around visual magnitude 13.0, and the other "inactive" or "off" state, where the brightness remains at about 15.0 magnitude (Hoffmeister et al 1985). These two states are thought to be the result of active and inactive mass-transfer rates from the secondary to the primary star (Hessman et al 2000).

Visual light curve of AM Herculis from the AAVSO International Database; August 1976 through March 2001.

Some of the short-term phenomena in the light curve of AM Her can be explained by the orbital motion of a binary system with a period of 3.1 hours. The 3.1 orbital period was discovered in AM Her from the eclipsing light-changes, the strongly variable linear and circular polarization, and the periodic radial velocity changes in the H and He lines (Hoffmeister et al 1985). Liller (1977) explains two kinds of optical variations associated with the orbital motion that are taking place in AM Her.

First, the red dwarf star is distorted into an egg-shaped figure by the attraction of its companion, toward which the long axis of the egg points. When we see the normal star broadside, it appears slightly brighter than when end on. Hence, as the entire system rotates, there are two long weak brightness maxima and two long shallow minima per period. Second, there can sometimes be observed brightness fluctuations due to the heating of the surface of the red secondary star by X-rays emitted by the collapsed star. This "hot spot" is periodically lost from view on the far side of the rotating normal star. Moreover, the short term light variations described earlier as "incessant flickering", occur due to the turbulent nature of the transfer of mass from the secondary to the white dwarf star (Hellier 2001).

The name polar was introduced by the Polish astronomers Krzeminski and Serkowski (1977) for AM Her and related objects, on account of the strong, and variable, linear and circular polarization in the light from these stars. Normal light is composed of electric vectors that are arranged randomly (always perpendicular to the direction of travel). Polarized light, in contrast to normal light, is composed of electric fields in a direction that is not random. If the electric vectors of a set of photons all point in one direction then the radiation is said to be linearly polarized.

AM Her displays both linear and circular polarized light for the following reason. The ionized material in an AM Her accretion stream does not simply follow a magnetic field line; instead it spirals around the field line (Figure 1). Consider looking at this field line side on (Figure 2). From this viewpoint, an electron spiraling around the field line will appear to be oscillating perpendicularly to the field line. The photons produced by this motion will always have an electric vector in the direction of this oscillation, and so the light is linearly polarized.

Now consider viewing the field line head on (Figure 3). The electron will appear to be circling, and so the electric vector of the emitted photons (which follows the electron's motion) will be continually rotating, tracing out a circle. Such light is said to be circularly polarized.

Thus, movement along the field lines beams linearly polarized light perpendicular to the field line and circularly polarized light along the field line. This type of emission is called cyclotron emission because spiraling electrons emit the photons that we see. Since the cyclotron emission can be as much as half the total light of AM Her stars, they are the most polarized objects in the sky (Hellier 2001).

Observing AM Her

It has only been a little more than two decades since AM Her was originally listed as the enigmatic "irregular" variable star in 1975. Since that time, astronomers have learned so much about AM Her that they have been able to create a good model for a new class of magnetic cataclysmic variable stars. Many of the questions were answered by correlating multi-wavelength data with optical data — AAVSO observations. Ground based observations like those that amateur astronomers can contribute have been, and will continue to be a critical link in understanding this star. A system as violent as AM Her cannot continue to exist in this state for long so it will be interesting to watch what happens to it over the next couple of decades. AM Her is a great target for interested CCD observers as it is usually found shining around visual magnitude 13.0 to 15.0. The AAVSO has a wide variety of charts available to help you locate AM Her in the sky. Also, a published AAVSO Monograph on AM Herculis contains a long-term light curve, which includes more than 10,000 visual observations of AM Her from 177 observers around the world.

These are all stars of B or A spectra concentrated in the galactic arms (Orion and Perseus), and they appear in OB associations or young galactic clusters. They therefore form a typical Population I.


X-ray Flare Stars

X-Ray Flare Stars produce flares generally lasting only a few minutes and sometimes less. At the time of maximum, the x-ray flux is one thousand to ten thousand times greater than the visible flux. The flares are often separated by an interval of several hours, of not several days.

An important point is that among the 30+ flare sources currently known most are concentrate towards the center of the Galaxy. In addition, several have been observed in globular clusters (NGC 851, NGC 6441, NGC 6624, NGC 6712).So it seems that these stars belong to Population II, which makes them very different from the red dwarfs we know as flare variables. Some of these x-ray sources can be identified, such as MXB 1827-05 now called MM Ser, or MXB 1733-44 (V926 Sco); they are optically variable and have a spectrum resembling that of dwarf novae.

It is thought that, here to, we have a binary containing a neutron star surrounded by a gaseous ring. The flares are probably produces by frequently repeated small explosions occurring in the ring.

Extra-galactic Monsters

Introduction


Seyfert Galaxies


BL Lacerta Objects


Quasars