ستارهٔ دوتایی به یک سامانهٔ ستارهای گفته میشود که در آن دو ستاره به هم نزدیک هستند و به دور مرکز سنگینی سراسری مشترک میان خود گردش میکنند. به ستارهٔ دیگر ستارهٔ ندیم یا ستارهٔ همدم نیز گفته میشود. بررسیهای جدید نشان میدهند که درصد زیادی از ستارگان بخشی از یک سامانهٔ حداقل دو ستارهای هستند. ستارگان دوتایی در اخترفیزیک بسیار مهم هستند زیرا ویژگیهای مدار آنها جرم و چگالی آن ستارگان را برای اخترشناسان مشخص میکند. جرم بسیاری از ستارگان تکی نیز از روی برونیابی جرم ستارگان دوتایی بدست میآید. ستارگان دوتایی واقعی با ستارگان دوتایی نوری یکی نیستند، تفاوت آنها در این است که ستارگان دوتایی نوری از زمین و از دیدگاه ما با چشم غیر مسلح نزدیک به یکدیگر یا گاهی به صورت یک ستاره دیده میشوند ولی آنها هیچ اثر گرانشی بر یکدیگر ندارند و فقط در راستای دید ناظر اینگونه دیده میشوند. ستارگان دوتایی از روی طیفسنجی هم شناخته میشوند. اگر مدار حرکت این ستارگان در راستای دید زمین باشد به آنها دوتایی گرفتی میگویند و هویت آنها از راه بهرهوری از پدیدهٔ گرفت تشخیص داده خواهدشد.
به سامانههایی که بیشتر از دو ستاره باهم مرکز سنگینی سراسری مشترک داشتهباشند ستارهٔ چندتایی میگویند که نمونهٔ آنها بسیار فراواناست. ستارههای دوتایی گاهی میتوانند بین یکدیگر جرم تبادل کنند و تکامل یابند. از معروفترین ستارگان دوتایی میتوان به الغول (ستارهٔ دوتایی گرفتی)، شباهنگ و ماکیان ایکس یک (که ندیم کوچکتر قویترین احتمال سیاهچاله است) اشاره کرد.
لفظ ستاره دوتایی از سال ۱۸۰۲ توسط سر ویلیام هرشل به کار رفت، در تعریف او آمدهاست: «یک ستاره دوتایی واقعی- متشکل از دو ستارهاست به طوری که به یکدیگر را جذب میکنند». دو ستاره نزدیک به هم بهطور تصادفیستاره دوتایی نوری نام میگیرند، مشهورترین ستاره دوتایی نوری ستاره زتا خرس بزرگ است که در صورت فلکی خرس بزرگ قرار دارد. دوتاییهای غیر واقعی را دوتایی نوری میگویند. با اختراع تلسکوپ ستارگان دوتایی بیشتری کشف شدند. هرشل، در ۱۷۸۰ میلادی، ۷۰۰ سامانه دوتایی، و حدود ۵۰ سامانه که به نظرش بیشتر از دو ستاره دارند ثبت نمود. یک سامانه دوتایی واقعی، دو ستارهاست که جاذبه گرانشی دارند. وقتی دو ستاره تفکیک میشوند که با بالابردن دقت تلسکوپها به اندازه کافی (در صورت لزوم از روشهای اینترفرمتریک استفاده میشود) دو ستاره کاملاً مجزا دیده شوند که به آنها دوتایی مرئی میگویند. نشان بروز دیگر تغییر طول موج بر اثر پدیده دوپلر است؛ که به این نوع ستارگان دوتایی دوتایی طیفی گفته میشود در این ستارهها خطوط طیفی بر اساس حرکت به سمت زمین یا مخالف آن به سمت ترتیب به آبی و قرمز حرکت میکنند، اگر خطوط به سمت آبی و سپس به سمت قرمز میل کنند نشان دهنده این است که این ستاره به دور مرکز ثقلی حرکت میکند، البته به شرط اینکه با دوره منظمی این اتفاق بیفتد اگر مدار آنها بسیار نزدیک به خط دید ما باشد، قدر ستاره بهطور منظمی کاهش و افزایش مییابد (البته با ستاره متغیر متفاوت است)، که به آنها دوتایی گرفتی میگویند بهترین مثال آن ستاره الغول است.
ستارگان دوتایی که هم مرئی و هم طیفی باشند، نادر هستند، و اگر یافت شوند اطلاعات مفیدی را در اختیار ستارهشناسان قرار میدهند چون ستارگان دوتایی بصری غالباً جدایی بزرگی در حالت واقعی دارند و دورههایشان بین دهها تا قرنهاست؛ ستارگان دوتایی که در مدار کوچکتری دارند، ستارگان دوتایی طیفی میشوند. به عبارت دیگر ستارگان دوتایی طیفی در مدار خود بسیار سریع حرکت میکنند زیرا مدار کوچکتری دارند بنابراین ستارگان دوتایی که هم طیفی باشند و هم مرئی باید بسیار به زمین نزدیک باشند تا دیده شوند.
اخترشناسان نوعی ستاره دوتایی کشف کردهاند که به آن دوتایی نجومسنجی میگویند این گونه ستارگان به ظاهر حول یک منطقه خالی چرخش میکنند که با بررسی بیشتر جرم زیادی در آن منطقه پیدا میشود که نامرئی است مانند سیاهچاله، ستاره نوترونی و... بعضی از دوتاییهای طیفی فقط یک خطوط طیفی متحرک دارند. بررسی بیشتر در طول موج پرتو ایکس در بعضی از این گروه ستارگان منبع بزرگی از پرتو ایکس در مرکز جرم را نشان میدهد. بهترین مثال آن ماکیان ایکس یک است، از روی جرم یافت شده که نه برابر جرم خورشید است و از روی حد تولمن-اوپنهایمر-ولکف (حداکثر مقدار جرم یک ستاره نوترونی از این مقدار کمتر است). این مرکز جرم بزرگترین نامزد برای سیاهچالهاست.
بر اساس روشهای رصد[ویرایش]
ستارگان دوتایی از لحاظ روشهای رصد به چهار گروه تقسیم میشوند. ستارگان دوتایی میتوانند در چند گروه نیز جای بگیرند اما این موضوع نادر است.
یک دوتایی مرئی مجموعهای از دو ستارهاست که مدارشان با راستای دید ما زاویه قابل توجه میسازد برای پیدا کردن این گونه ستارگان باید از تلسکوپ با قدرت تفکیک کافی استفاده کرد و هرچه دهانه تلسکوپ بزرگتر باشد قدرت تفکیک دو ستاره از هم بالا میرود.
به ستارهٔ درخشانتر در ستارهٔ دوتایی مرئی ستارهٔ اولیه(یا اصلی) و به ستارهٔ کمنورتر ستارهٔ ثانویه(یا همدم) گفته میشود. اگر دو ستاره قدر ظاهری یکسان داشته باشند به ستارهای اصلی گفته میشود که زودتر کشف شده باشد.
گاهی اوقات، تنها گواه دوتایی بودن یک ستاره با بررسی اثر داپلر بر روی نور ساطع شده از آناست. در چنین موردی، اگر ستارهٔ دوتایی متشکل از یک جفت ستاره باشد در آن از خطوط طیفی در نور ساطع شده از هر ستاره، نخست به سمت آبی کشانده میشود، و پس از آن، به سمت رنگ قرمز متمایل میشود، که همان حرکت به سوی ما، و سپس دور شدن از ما، در طول حرکت خود به دور مرکز سنگینی سراسری مشترک خود، در طول دور زدن مدار خویش است.
یک دوتایی گرفتی دو ستارهاست که مدار آنها بسیار نزدیک به راستای دید ناظر باشد در این صورت در طول دوره دو بار گرفت رخ خواهد داد. حال اگر دوتایی گرفتی طیفی هم باشد و اختلافمنظر سامانه نیز در دست باشد، چنین سامانهای برای تجزیه وتحلیل اخترشناسی بسیار پرارزش خواهد بود.الغول بهترین نمونهٔ تاکنون شناخته شدهٔ دوتایی گرفتیاست.
در دههٔ گذشته، اندازهگیری پارامترهای اساسی دوتاییهای گرفتی فراکهکشانی با تلسکوپهای سری ۸ متری امکانپذیر شدهاست. این امکان باعث شده که از آنها برای بهطور مستقیم اندازهگیری فاصلهٔ کهکشانهای دور، به صورت بسیار دقیق تر استفاده شود، روشی که از به کارگیری شمع استاندارد بسیار دقیق تراست.
بر اساس شکل سیستم[ویرایش]
تکامل ستارگان دوتایی[ویرایش]
انتقال جرم و یکپارچگی[ویرایش]
Runaways and novae[ویرایش]
کاربرد در اخترفیزیک[ویرایش]
نتیجهگیری بررسیهای بیشتر مرکز اخترفیزیک هاروارد-اسمیتسونین (۲۰۰۶) بر پایهٔ تحقیقات چارلز لادا برآورد کرده که؛ مغایر با فرض کنونی حدود یک سوم از سیستمهای ستارهای در کهکشان راه شیری یا دوتایی (باینری) یا چندستارهای هستند، و دو سوم باقی مانده را ستارگان منفرد تشکیل میدهند. از آنجا که سیارهها احتمالاً راحتتر در پیرامون ستارگان تنها شکل میگیرند، سیارهها نیز ممکن است متداولتر از آنچه قبلاً تصور شده باشند. این واقعیت که؛ ستارگان بزرگ و درخشانی مانند خورشید بیشتر به صورت دوتایی یا چندستارهای هستند، این تصور را ایجاد کردهبود که بیشتر ستارگان جهان چندستارهای هستند، در حالی که تحقیقات جدید پیرامون ستارگان کوچکتر و کمنورتر و کوتولهها نشان میدهد که درصد آنها بیش از ستارگان وابسته به سیستمهای دوتایی و چندستارهای است.
سیارات درکنار ستارگان دوتایی[ویرایش]
مثالهایی از ستارگان چندتایی[ویرایش]
پیوند به بیرون[ویرایش]
A binary star is a star system consisting of two stars orbiting around their common barycenter. Systems of two or more stars are called multiple star systems. These systems, especially when more distant, often appear to the unaided eye as a single point of light, and are then revealed as multiple by other means. Research over the last two centuries suggests that half or more of visible stars are part of multiple star systems.
The term double star is often used synonymously with binary star; however, double star can also mean optical double star. Optical doubles are so called because the two stars appear close together in the sky as seen from the Earth; they are almost on the same line of sight. Nevertheless, their "doubleness" depends only on this optical effect; the stars themselves are distant from one another and share no physical connection. A double star can be revealed as optical by means of differences in their parallax measurements, proper motions, or radial velocities. Most known double stars have not been studied adequately to determine whether they are optical doubles or doubles physically bound through gravitation into a multiple star system.
Binary star systems are very important in astrophysics because calculations of their orbits allow the masses of their component stars to be directly determined, which in turn allows other stellar parameters, such as radius and density, to be indirectly estimated. This also determines an empirical mass-luminosity relationship (MLR) from which the masses of single stars can be estimated.
Binary stars are often detected optically, in which case they are called visual binaries. Many visual binaries have long orbital periods of several centuries or millennia and therefore have orbits which are uncertain or poorly known. They may also be detected by indirect techniques, such as spectroscopy (spectroscopic binaries) or astrometry (astrometric binaries). If a binary star happens to orbit in a plane along our line of sight, its components will eclipse and transit each other; these pairs are called eclipsing binaries, or, as they are detected by their changes in brightness during eclipses and transits, photometric binaries.
If components in binary star systems are close enough they can gravitationally distort their mutual outer stellar atmospheres. In some cases, these close binary systems can exchange mass, which may bring their evolution to stages that single stars cannot attain. Examples of binaries are Sirius, and Cygnus X-1 (Cygnus X-1 being a well-known black hole). Binary stars are also common as the nuclei of many planetary nebulae, and are the progenitors of both novae and type Ia supernovae.
By the modern definition, the term binary star is generally restricted to pairs of stars which revolve around a common center of mass. Binary stars which can be resolved with a telescope or interferometric methods are known as visual binaries. For most of the known visual binary stars one whole revolution has not been observed yet, they are observed to have travelled along a curved path or a partial arc.
The more general term double star is used for pairs of stars which are seen to be close together in the sky. This distinction is rarely made in languages other than English. Double stars may be binary systems or may be merely two stars that appear to be close together in the sky but have vastly different true distances from the Sun. The latter are termed optical doubles or optical pairs.
Since the invention of the telescope, many pairs of double stars have been found. Early examples include Mizar and Acrux. Mizar, in the Big Dipper (Ursa Major), was observed to be double by Giovanni Battista Riccioli in 1650 (and probably earlier by Benedetto Castelli and Galileo). The bright southern star Acrux, in the Southern Cross, was discovered to be double by Father Fontenay in 1685.
John Michell was the first to suggest that double stars might be physically attached to each other when he argued in 1767 that the probability that a double star was due to a chance alignment was small. William Herschel began observing double stars in 1779 and soon thereafter published catalogs of about 700 double stars. By 1803, he had observed changes in the relative positions in a number of double stars over the course of 25 years, and concluded that they must be binary systems; the first orbit of a binary star, however, was not computed until 1827, when Félix Savary computed the orbit of Xi Ursae Majoris. Since this time, many more double stars have been catalogued and measured. The Washington Double Star Catalog, a database of visual double stars compiled by the United States Naval Observatory, contains over 100,000 pairs of double stars, including optical doubles as well as binary stars. Orbits are known for only a few thousand of these double stars, and most have not been ascertained to be either true binaries or optical double stars. This can be determined by observing the relative motion of the pairs. If the motion is part of an orbit, or if the stars have similar radial velocities and the difference in their proper motions is small compared to their common proper motion, the pair is probably physical. One of the tasks that remains for visual observers of double stars is to obtain sufficient observations to prove or disprove gravitational connection.
Methods of observation
Binary stars are classified into four types according to the way in which they are observed: visually, by observation; spectroscopically, by periodic changes in spectral lines; photometrically, by changes in brightness caused by an eclipse; or astrometrically, by measuring a deviation in a star's position caused by an unseen companion. Any binary star can belong to several of these classes; for example, several spectroscopic binaries are also eclipsing binaries.
A visual binary star is a binary star for which the angular separation between the two components is great enough to permit them to be observed as a double star in a telescope, or even high-powered binoculars. The angular resolution of the telescope is an important factor in the detection of visual binaries, and as better angular resolutions are applied to binary star observations, an increasing number of visual binaries will be detected. The relative brightness of the two stars is also an important factor, as glare from a bright star may make it difficult to detect the presence of a fainter component.
The brighter star of a visual binary is the primary star, and the dimmer is considered the secondary. In some publications (especially older ones), a faint secondary is called the comes (plural comites; companion). If the stars are the same brightness, the discoverer designation for the primary is customarily accepted.
The position angle of the secondary with respect to the primary is measured, together with the angular distance between the two stars. The time of observation is also recorded. After a sufficient number of observations are recorded over a period of time, they are plotted in polar coordinates with the primary star at the origin, and the most probable ellipse is drawn through these points such that the Keplerian law of areas is satisfied. This ellipse is known as the apparent ellipse, and is the projection of the actual elliptical orbit of the secondary with respect to the primary on the plane of the sky. From this projected ellipse the complete elements of the orbit may be computed, where the semi-major axis can only be expressed in angular units unless the stellar parallax, and hence the distance, of the system is known.
Sometimes, the only evidence of a binary star comes from the Doppler effect on its emitted light. In these cases, the binary consists of a pair of stars where the spectral lines in the light emitted from each star shifts first towards the blue, then towards the red, as each moves first towards us, and then away from us, during its motion about their common center of mass, with the period of their common orbit.
In these systems, the separation between the stars is usually very small, and the orbital velocity very high. Unless the plane of the orbit happens to be perpendicular to the line of sight, the orbital velocities will have components in the line of sight and the observed radial velocity of the system will vary periodically. Since radial velocity can be measured with a spectrometer by observing the Doppler shift of the stars' spectral lines, the binaries detected in this manner are known as spectroscopic binaries. Most of these cannot be resolved as a visual binary, even with telescopes of the highest existing resolving power.
In some spectroscopic binaries, spectral lines from both stars are visible and the lines are alternately double and single. Such a system is known as a double-lined spectroscopic binary (often denoted "SB2"). In other systems, the spectrum of only one of the stars is seen and the lines in the spectrum shift periodically towards the blue, then towards red and back again. Such stars are known as single-lined spectroscopic binaries ("SB1").
The orbit of a spectroscopic binary is determined by making a long series of observations of the radial velocity of one or both components of the system. The observations are plotted against time, and from the resulting curve a period is determined. If the orbit is circular then the curve will be a sine curve. If the orbit is elliptical, the shape of the curve will depend on the eccentricity of the ellipse and the orientation of the major axis with reference to the line of sight.
It is impossible to determine individually the semi-major axis a and the inclination of the orbit plane i. However, the product of the semi-major axis and the sine of the inclination (i.e. a sin i) may be determined directly in linear units (e.g. kilometres). If either a or i can be determined by other means, as in the case of eclipsing binaries, a complete solution for the orbit can be found.
Binary stars that are both visual and spectroscopic binaries are rare, and are a valuable source of information when found. About 40 are known. Visual binary stars often have large true separations, with periods measured in decades to centuries; consequently, they usually have orbital speeds too small to be measured spectroscopically. Conversely, spectroscopic binary stars move fast in their orbits because they are close together, usually too close to be detected as visual binaries. Binaries that are found to be both visual and spectroscopic thus must be relatively close to Earth.
An eclipsing binary star is a binary star system in which the orbit plane of the two stars lies so nearly in the line of sight of the observer that the components undergo mutual eclipses. In the case where the binary is also a spectroscopic binary and the parallax of the system is known, the binary is quite valuable for stellar analysis. Algol, a triple star system in the constellation Perseus, contains the best-known example of an eclipsing binary.
Eclipsing binaries are variable stars, not because the light of the individual components vary but because of the eclipses. The light curve of an eclipsing binary is characterized by periods of practically constant light, with periodic drops in intensity when one star passes in front of the other. The brightness may drop twice during the orbit, once when the secondary passes in front of the primary and once when the primary passes in front of the secondary. The deeper of the two eclipses is called the primary regardless of which star is being occulted, and if a shallow second eclipse also occurs it is called the secondary eclipse. The size of the brightness drops depends on the relative brightness of the two stars, the proportion of the occulted star that is hidden, and the surface brightness (i.e. effective temperature) of the stars. Typically the occultation of the hotter star causes the primary eclipse.
An eclipsing binaries' period of orbit may be determined from a study of its light curve, and the relative sizes of the individual stars can be determined in terms of the radius of the orbit, by observing how quickly the brightness changes as the disc of the nearest star slides over the disc of the other star. If it is also a spectroscopic binary, the orbital elements can also be determined, and the mass of the stars can be determined relatively easily, which means that the relative densities of the stars can be determined in this case.
Since about 1995, measurement of extragalactic eclipsing binaries' fundamental parameters has become possible with 8-meter class telescopes. This makes it feasible to use them to directly measure the distances to external galaxies, a process that is more accurate than using standard candles. By 2006, they had been used to give direct distance estimates to the LMC, SMC, Andromeda Galaxy, and Triangulum Galaxy. Eclipsing binaries offer a direct method to gauge the distance to galaxies to an improved 5% level of accuracy.
Non-eclipsing binaries that can be detected through photometry
Nearby non-eclipsing binaries can also be photometrically detected by observing how the stars affect each other in three ways. The first is by observing extra light which the stars reflect from their companion. Second is by observing ellipsoidal light variations which are caused by deformation of the star's shape by their companions. The third method is by looking at how relativistic beaming affects the apparent magnitude of the stars. Detecting binaries with these methods requires accurate photometry.
Astronomers have discovered some stars that seemingly orbit around an empty space. Astrometric binaries are relatively nearby stars which can be seen to wobble around a point in space, with no visible companion. The same mathematics used for ordinary binaries can be applied to infer the mass of the missing companion. The companion could be very dim, so that it is currently undetectable or masked by the glare of its primary, or it could be an object that emits little or no electromagnetic radiation, for example a neutron star.
The visible star's position is carefully measured and detected to vary, due to the gravitational influence from its counterpart. The position of the star is repeatedly measured relative to more distant stars, and then checked for periodic shifts in position. Typically this type of measurement can only be performed on nearby stars, such as those within 10 parsecs. Nearby stars often have a relatively high proper motion, so astrometric binaries will appear to follow a wobbly path across the sky.
If the companion is sufficiently massive to cause an observable shift in position of the star, then its presence can be deduced. From precise astrometric measurements of the movement of the visible star over a sufficiently long period of time, information about the mass of the companion and its orbital period can be determined. Even though the companion is not visible, the characteristics of the system can be determined from the observations using Kepler's laws.
This method of detecting binaries is also used to locate extrasolar planets orbiting a star. However, the requirements to perform this measurement are very exacting, due to the great difference in the mass ratio, and the typically long period of the planet's orbit. Detection of position shifts of a star is a very exacting science, and it is difficult to achieve the necessary precision. Space telescopes can avoid the blurring effect of Earth's atmosphere, resulting in more precise resolution.
Configuration of the system
Another classification is based on the distance between the stars, relative to their sizes:
Detached binaries are binary stars where each component is within its Roche lobe, i.e. the area where the gravitational pull of the star itself is larger than that of the other component. The stars have no major effect on each other, and essentially evolve separately. Most binaries belong to this class.
Semidetached binary stars are binary stars where one of the components fills the binary star's Roche lobe and the other does not. Gas from the surface of the Roche-lobe-filling component (donor) is transferred to the other, accreting star. The mass transfer dominates the evolution of the system. In many cases, the inflowing gas forms an accretion disc around the accretor.
A contact binary is a type of binary star in which both components of the binary fill their Roche lobes. The uppermost part of the stellar atmospheres forms a common envelope that surrounds both stars. As the friction of the envelope brakes the orbital motion, the stars may eventually merge. W Ursae Majoris is an example.
Cataclysmic variables and X-ray binaries
When a binary system contains a compact object such as a white dwarf, neutron star or black hole, gas from the other (donor) star can accrete onto the compact object. This releases gravitational potential energy, causing the gas to become hotter and emit radiation. Cataclysmic variable stars, where the compact object is a white dwarf, are examples of such systems. In X-ray binaries, the compact object can be either a neutron star or a black hole. These binaries are classified as low-mass or high-mass according to the mass of the donor star. High-mass X-ray binaries contain a young, early-type, high-mass donor star which transfers mass by its stellar wind, while low-mass X-ray binaries are semidetached binaries in which gas from a late-type donor star or a white dwarf overflows the Roche lobe and falls towards the neutron star or black hole. Probably the best known example of an X-ray binary is the high-mass X-ray binary Cygnus X-1. In Cygnus X-1, the mass of the unseen companion is estimated to be about nine times that of the Sun, far exceeding the Tolman–Oppenheimer–Volkoff limit for the maximum theoretical mass of a neutron star. It is therefore believed to be a black hole; it was the first object for which this was widely believed.
Variations in period
The Applegate mechanism explains long term orbital period variations seen in certain eclipsing binaries. As a main-sequence star goes through an activity cycle, the outer layers of the star are subject to a magnetic torque changing the distribution of angular momentum, resulting in a change in the star's oblateness. The orbit of the stars in the binary pair is gravitationally coupled to their shape changes, so that the period shows modulations (typically on the order of ∆P/P ∼ 10−5) on the same time scale as the activity cycles (typically on the order of decades).
Another phenomenon observed in some Algol binaries has been monotonic period increases. This is quite distinct from the far more common observations of alternating period increases and decreases explained by the Applegate mechanism. Monotonic period increases have been attributed to mass transfer, usually (but not always) from the less massive to the more massive star
A and B
The components of binary stars are denoted by the suffixes A and B appended to the system's designation, A denoting the primary and B the secondary. The suffix AB may be used to denote the pair (for example, the binary star α Centauri AB consists of the stars α Centauri A and α Centauri B.) Additional letters, such as C, D, etc., may be used for systems with more than two stars. In cases where the binary star has a Bayer designation and is widely separated, it is possible that the members of the pair will be designated with superscripts; an example is Zeta Reticuli, whose components are ζ1 Reticuli and ζ2 Reticuli.
Double stars are also designated by an abbreviation giving the discoverer together with an index number. α Centauri, for example, was found to be double by Father Richaud in 1689, and so is designated RHD 1. These discoverer codes can be found in the Washington Double Star Catalog.
Hot and cold
The components of a binary star system may be designated by their relative temperatures as the hot companion and cool companion.
While it is not impossible that some binaries might be created through gravitational capture between two single stars, given the very low likelihood of such an event (three objects being actually required, as conservation of energy rules out a single gravitating body capturing another) and the high number of binaries currently in existence, this cannot be the primary formation process. The observation of binaries consisting of stars not yet on the main sequence supports the theory that binaries develop during star formation. Fragmentation of the molecular cloud during the formation of protostars is an acceptable explanation for the formation of a binary or multiple star system.
The outcome of the three-body problem, in which the three stars are of comparable mass, is that eventually one of the three stars will be ejected from the system and, assuming no significant further perturbations, the remaining two will form a stable binary system.
Mass transfer and accretion
As a main-sequence star increases in size during its evolution, it may at some point exceed its Roche lobe, meaning that some of its matter ventures into a region where the gravitational pull of its companion star is larger than its own. The result is that matter will transfer from one star to another through a process known as Roche lobe overflow (RLOF), either being absorbed by direct impact or through an accretion disc. The mathematical point through which this transfer happens is called the first Lagrangian point. It is not uncommon that the accretion disc is the brightest (and thus sometimes the only visible) element of a binary star.
If a star grows outside of its Roche lobe too fast for all abundant matter to be transferred to the other component, it is also possible that matter will leave the system through other Lagrange points or as stellar wind, thus being effectively lost to both components. Since the evolution of a star is determined by its mass, the process influences the evolution of both companions, and creates stages that cannot be attained by single stars.
Studies of the eclipsing ternary Algol led to the Algol paradox in the theory of stellar evolution: although components of a binary star form at the same time, and massive stars evolve much faster than the less massive ones, it was observed that the more massive component Algol A is still in the main sequence, while the less massive Algol B is a subgiant at a later evolutionary stage. The paradox can be solved by mass transfer: when the more massive star became a subgiant, it filled its Roche lobe, and most of the mass was transferred to the other star, which is still in the main sequence. In some binaries similar to Algol, a gas flow can actually be seen.
Runaways and novae
It is also possible for widely separated binaries to lose gravitational contact with each other during their lifetime, as a result of external perturbations. The components will then move on to evolve as single stars. A close encounter between two binary systems can also result in the gravitational disruption of both systems, with some of the stars being ejected at high velocities, leading to runaway stars.
If a white dwarf has a close companion star that overflows its Roche lobe, the white dwarf will steadily accrete gases from the star's outer atmosphere. These are compacted on the white dwarf's surface by its intense gravity, compressed and heated to very high temperatures as additional material is drawn in. The white dwarf consists of degenerate matter and so is largely unresponsive to heat, while the accreted hydrogen is not. Hydrogen fusion can occur in a stable manner on the surface through the CNO cycle, causing the enormous amount of energy liberated by this process to blow the remaining gases away from the white dwarf's surface. The result is an extremely bright outburst of light, known as a nova.
In extreme cases this event can cause the white dwarf to exceed the Chandrasekhar limit and trigger a supernova that destroys the entire star, another possible cause for runaways. An example of such an event is the supernova SN 1572, which was observed by Tycho Brahe. The Hubble Space Telescope recently took a picture of the remnants of this event.
Binaries provide the best method for astronomers to determine the mass of a distant star. The gravitational pull between them causes them to orbit around their common center of mass. From the orbital pattern of a visual binary, or the time variation of the spectrum of a spectroscopic binary, the mass of its stars can be determined, for example with the binary mass function. In this way, the relation between a star's appearance (temperature and radius) and its mass can be found, which allows for the determination of the mass of non-binaries.
Because a large proportion of stars exist in binary systems, binaries are particularly important to our understanding of the processes by which stars form. In particular, the period and masses of the binary tell us about the amount of angular momentum in the system. Because this is a conserved quantity in physics, binaries give us important clues about the conditions under which the stars were formed.
Calculating the center of mass in binary stars
In a simple binary case, r1, the distance from the center of the first star to the center of mass or barycenter, is given by:
If a is taken to be the semi-major axis of the orbit of one body around the other, then r1 will be the semimajor axis of the first body's orbit around the center of mass or barycenter, and r2 = a – r1 will be the semimajor axis of the second body's orbit. When the center of mass is located within the more massive body, that body will appear to wobble rather than following a discernible orbit.
Center of mass animations
Images are representative, not simulated. The position of the red cross indicates the center of mass of the system.
It is estimated that approximately one third of the star systems in the Milky Way are binary or multiple, with the remaining two thirds being single stars. The overall multiplicity frequency of ordinary stars is a monotonically increasing function of stellar mass. That is, the likelihood of being in a binary or a multi-star system steadily increases as the mass of the components increase.
There is a direct correlation between the period of revolution of a binary star and the eccentricity of its orbit, with systems of short period having smaller eccentricity. Binary stars may be found with any conceivable separation, from pairs orbiting so closely that they are practically in contact with each other, to pairs so distantly separated that their connection is indicated only by their common proper motion through space. Among gravitationally bound binary star systems, there exists a so-called log normal distribution of periods, with the majority of these systems orbiting with a period of about 100 years. This is supporting evidence for the theory that binary systems are formed during star formation.
In pairs where the two stars are of equal brightness, they are also of the same spectral type. In systems where the brightnesses are different, the fainter star is bluer if the brighter star is a giant star, and redder if the brighter star belongs to the main sequence.
The mass of a star can be directly determined only from its gravitational attraction. Apart from the Sun and stars which act as gravitational lenses, this can be done only in binary and multiple star systems, making the binary stars an important class of stars. In the case of a visual binary star, after the orbit and the stellar parallax of the system has been determined, the combined mass of the two stars may be obtained by a direct application of the Keplerian harmonic law.
Unfortunately, it is impossible to obtain the complete orbit of a spectroscopic binary unless it is also a visual or an eclipsing binary, so from these objects only a determination of the joint product of mass and the sine of the angle of inclination relative to the line of sight is possible. In the case of eclipsing binaries which are also spectroscopic binaries, it is possible to find a complete solution for the specifications (mass, density, size, luminosity, and approximate shape) of both members of the system.
While a number of binary star systems have been found to harbor extrasolar planets, such systems are comparatively rare compared to single star systems. Observations by the Kepler space telescope have shown that most single stars of the same type as the Sun have plenty of planets, but only one-third of binary stars do. According to theoretical simulations, even widely separated binary stars often disrupt the discs of rocky grains from which protoplanets form. On the other hand, other simulations suggest that the presence of a binary companion can actually improve the rate of planet formation within stable orbital zones by "stirring up" the protoplanetary disk, increasing the accretion rate of the protoplanets within.
Detecting planets in multiple star systems introduces additional technical difficulties, which may be why they are only rarely found. Examples include the white dwarf-pulsar binary PSR B1620-26, the subgiant-red dwarf binary Gamma Cephei, and the white dwarf-red dwarf binary NN Serpentis; among others.
A study of fourteen previously known planetary systems found three of these systems to be binary systems. All planets were found to be in S-type orbits around the primary star. In these three cases the secondary star was much dimmer than the primary and so was not previously detected. This discovery resulted in a recalculation of parameters for both the planet and the primary star.
Science fiction has often featured planets of binary or ternary stars as a setting, for example George Lucas' Tatooine from Star Wars, and one notable story, "Nightfall", even takes this to a six-star system. In reality, some orbital ranges are impossible for dynamical reasons (the planet would be expelled from its orbit relatively quickly, being either ejected from the system altogether or transferred to a more inner or outer orbital range), whilst other orbits present serious challenges for eventual biospheres because of likely extreme variations in surface temperature during different parts of the orbit. Planets that orbit just one star in a binary system are said to have "S-type" orbits, whereas those that orbit around both stars have "P-type" or "circumbinary" orbits. It is estimated that 50–60% of binary systems are capable of supporting habitable terrestrial planets within stable orbital ranges.
The large distance between the components, as well as their difference in color, make Albireo one of the easiest observable visual binaries. The brightest member, which is the third-brightest star in the constellation Cygnus, is actually a close binary itself. Also in the Cygnus constellation is Cygnus X-1, an X-ray source considered to be a black hole. It is a high-mass X-ray binary, with the optical counterpart being a variable star. Sirius is another binary and the brightest star in the night time sky, with a visual apparent magnitude of −1.46. It is located in the constellation Canis Major. In 1844 Friedrich Bessel deduced that Sirius was a binary. In 1862 Alvan Graham Clark discovered the companion (Sirius B; the visible star is Sirius A). In 1915 astronomers at the Mount Wilson Observatory determined that Sirius B was a white dwarf, the first to be discovered. In 2005, using the Hubble Space Telescope, astronomers determined Sirius B to be 12,000 km (7,456 mi) in diameter, with a mass that is 98% of the Sun.
An example of an eclipsing binary is Epsilon Aurigae in the constellation Auriga. The visible component belongs to the spectral class F0, the other (eclipsing) component is not visible. The last such eclipse occurred from 2009–2011, and it is hoped that the extensive observations that will likely be carried out may yield further insights into the nature of this system. Another eclipsing binary is Beta Lyrae, which is a semidetached binary star system in the constellation of Lyra.
Other interesting binaries include 61 Cygni (a binary in the constellation Cygnus, composed of two K class (orange) main-sequence stars, 61 Cygni A and 61 Cygni B, which is known for its large proper motion), Procyon (the brightest star in the constellation Canis Minor and the eighth-brightest star in the night time sky, which is a binary consisting of the main star with a faint white dwarf companion), SS Lacertae (an eclipsing binary which stopped eclipsing), V907 Sco (an eclipsing binary which stopped, restarted, then stopped again) and BG Geminorum (an eclipsing binary which is thought to contain a black hole with a K0 star in orbit around it), 2MASS J18082002-5104378 (a binary in the "thin disk" of the Milky Way, and containing one of the oldest known stars).
Multiple star examples
Systems with more than two stars are termed multiple stars. Algol is the most noted ternary (long thought to be a binary), located in the constellation Perseus. Two components of the system eclipse each other, the variation in the intensity of Algol first being recorded in 1670 by Geminiano Montanari. The name Algol means "demon star" (from Arabic: الغول al-ghūl), which was probably given due to its peculiar behavior. Another visible ternary is Alpha Centauri, in the southern constellation of Centaurus, which contains the fourth-brightest star in the night sky, with an apparent visual magnitude of −0.01. This system also underscores the fact that no search for habitable planets is complete if binaries are discounted. Alpha Centauri A and B have an 11 AU distance at closest approach, and both should have stable habitable zones.
There are also examples of systems beyond ternaries: Castor is a sextuple star system, which is the second-brightest star in the constellation Gemini and one of the brightest stars in the nighttime sky. Astronomically, Castor was discovered to be a visual binary in 1719. Each of the components of Castor is itself a spectroscopic binary. Castor also has a faint and widely separated companion, which is also a spectroscopic binary. The Alcor–Mizar visual binary in Ursa Majoris also consists of six stars, four comprising Mizar and two comprising Alcor.
Notes and references