True or False a Strong Stellar Wind Causes a Red Giant to Continually Lose Mass

Stellar Wind

Stars, Massive

Steven N. Shore , in Encyclopedia of Physical Science and Technology (Third Edition), 2003

III.B Main Sequence Stages

During most of their lifetimes, massive stars are dominated by their stellar winds. Their masses are never constant, but are constantly decreasing as their luminosities change. On the main sequence they burn hydrogen into helium via the CNO process, during which phase they possess massive convective cores and radiative envelopes. The stellar luminosity of the mast massive main sequence stars is very close to the Eddington limit for most of their lives. Maeder (1987) gives the following mass–luminosity relations for the massive stars on the main sequence:

(6) $log\frac{\text{L}}{{\text{L}}_{\odot }}=2.55log\frac{\text{M}}{{\text{M}}_{\odot }}+1.27$

for 15   ≤ M/M _⊙  ≤   40 and

(7) $log\frac{\text{L}}{{\text{L}}_{\odot }}=1.84log\frac{\text{M}}{{\text{M}}_{\odot }}+2.42$

for 40   ≤ M/M_⊙  ≤   120. It is possible that the most massive stars, above about 60 M _⊙, do not have a real main sequence stage in the sense of being globally static. Their surfaces may be merely opaque layers in their wind. As a result, the atmosphere is likely dynamic and more extended than a stationary layer. This, if correct, alters the age determination from fitting the main sequence isochrones to clusters that are computed with conventional assumptions.

Stars more massive than about 25 M _⊙ lose as much as, or more than, 10% of their mass while still in the hydrogen core burning phase of evolution. There is a rapid increase in this rate for the stars more massive than 25 M _⊙, reaching upward of 30% for 80 M _⊙ stars. The main sequence lifetime for these stars is more difficult to assess than for lower mass objects. Because their mass continually decreases even during the hydrogen core burning stage of evolution, their lifetimes are longer than would be anticipated from constant-mass models:

(8) ${\text{t}}_{\text{MS}}=1.3\times {10}^8{\left(\frac{\text{M}}{{\text{M}}_{\odot }}\right)}^{-0.86}\text{yr}$

for stars with masses 15   ≤ M/M _⊙  ≤   60. The greatest uncertainty in these figures is the rate of mass loss while the star is on the main sequence. This depends critically on the prescription used for calculating the relation between the stellar luminosity and the mass loss in the stellar wind.

Winds are observed for luminous stars. The question is not whether they occur, but what drives them and how the driving depends on the stellar properties, such as mass, radius, and luminosity. If driven by turbulence or some other form of direct mechanical input, the dependence on the luminosity is through the surface temperature and radius of the star. On the other hand, for radiatively driven winds, the driving depends on the envelope opacity and consequently the metal abundances. Presently, there is considerable uncertainty associated with these choices and the reader is referred to Maeder and Chiosi (1994) for a tabulated listing of currently proposed theoretical and empirical laws for M as a function of stellar parameters.

The main sequence lifetime is also affected by core convection. There are still uncertainties in the physics of stellar convective energy transport, particularly concerning the overshooting of convective cells at the boundary of the CNO-processed core and overlying radiative envelope. Mixing of envelope material, which is hydrogen rich, into the helium core extends the lifetime of the star in the hydrogen core burning stage and increases the core mass relative to the envelope. It also extends the period of subsequent helium core burning and thus alters the entire subsequent lifetime of the star. While population statistics of stellar distributions near the main sequence on the Hertzsprung–Russell diagram (HRD) seem to require some mixing during this early stage of evolution, the mechanism for such mixing is an open problem.

As helium builds in the core, the star begins to evolve toward the red giant branch, with increasing luminosity and larger radius. Here again, the effects of mass loss play an important role. As the surface gravity drops and the luminous output of the nuclear source increases, the stellar wind becomes more powerful. The combined effect of core contraction and increased mass loss produces nearly constant luminosity evolution across the HRD. The star's trajectory is then cutting across isochrones computed for constant mass models with more than 30% of the mass lost by the time of carbon core ignition. The core of a massive star evolves almost independently from the mass of the envelope, so the nucleosynthetic details are less sensitive to mass loss prescription than are the surface properties.

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Supernovae

David Branch , in Encyclopedia of Physical Science and Technology (Third Edition), 2003

IV.C Circumstellar Interaction

The progenitor of a core-collapse supernova loses mass by means of a stellar wind. For a steady-state wind that has a constant velocity and a constant mass-loss rate, the density of the circumstellar matter is proportional to the mass-loss rate, inversely proportional to the wind velocity, and inversely proportional to the square of the distance from the star. The wind velocity ordinarily is comparable to the escape velocity from the photosphere of the star—of the order of 10  km/sec for a red supergiant and 1000   km/sec for a blue one—so the circumstellar matter of a red supergiant is much more dense than that of a blue one. When a core-collapse supernova occurs, both the ejected matter and the radiated photons interact with the circumstellar matter. The violent collision between the high-velocity ejecta and the circumstellar matter is called a hydrodynamic interaction. The interaction between the supernova photons and the circumstellar matter is called a radiative interaction. These interactions can have observable consequences in various parts of the electromagnetic spectrum.

Hydrodynamic circumstellar interaction has been detected most often by observations at radio wavelengths. The spectrum of the radio emission indicates that it is synchrotron radiation from electrons that are accelerated to relativistic velocities in a very hot (up to 10⁹  K) interaction region. The higher the density of the circumstellar matter, the longer the radio "light curve" takes to rise to its peak, because initially the circumstellar matter outside the interaction region, which has been ionized by the X-ray and ultraviolet flash, absorbs the radio photons by free–free processes. Thermal X-rays from the hot interaction region also have been detected in a smaller number of cases. The fraction of a supernova's total luminosity that is emitted in the radio and X-ray bands ordinarily is very small, but observation of the radio and X-ray emission combined with modeling of the interaction provides valuable information on the mass-loss rate of the progenitor star. Rates in the range of 10⁻⁶ to 10⁻⁴ solar masses per year have been inferred. Optical light radiated from the interaction region is responsible for the slow decay rate of the light curves of SNe IIn. Unlike most supernovae, which seldom are followed observationally for more than a year after explosion, some circumstellar-interacting supernovae are observed decades after explosion, at optical, radio, and X-ray wavelengths.

Radiative interactions are manifested most clearly by relatively narrow emission and absorption lines that appear in optical spectra (which cause the supernova to be classified Type IIn) and, especially, in ultraviolet spectra [which can only be obtained from above Earth's atmosphere with instruments such as the Hubble Space Telescope (HST); Fig. 10]. The widths of the circumstellar lines provide information on the wind velocity of the progenitor, and analysis of the line strengths provides information on the relative abundances of the elements in the wind. The ionization state of the circumstellar matter also provides information on the amount of ionizing radiation that was emitted by the explosion, something that is not observable directly.

FIGURE 10. Ultraviolet spectra of the Type IIn SN 1998S obtained with the Hubble Space Telescope. [Courtesy of P. Challis, Harvard–Smithsonian Center for Astrophysics and Supernova INtensive Study (SINS) team.]

Some supernovae emit excess infrared radiation, well beyond the amounts expected from their photospheres, by thermal radiation from cool (1000   K), small (10⁻⁵  cm) solid particles (dust grains) in the circumstellar medium. The grains are heated by absorbing optical and ultraviolet radiation from the supernova and respond by emitting in the infrared. Infrared observations provide information on the nature and spatial distribution of the dust grains.

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Interstellar Matter

Donald G. York , in Encyclopedia of Physical Science and Technology (Third Edition), 2003

VI.D Evolution

The evolution of the intercloud medium depends on the injection of ionization energy through supernova blast waves, UV photons, and stellar winds. A single supernova may keep a region of 100-pc diameter ionized for 10 ⁶ years because of the small cooling rate of such hot, low-density gas. Ionizing photons from O stars in a region free of dense clouds may ionize a region as large as 30–100   pc in diameter for 10⁶ years before all the stellar nuclear fuel is exhausted. Thus in star-forming regions of galaxies with low ambient densities and with supernova rates of 1 per 10⁶ years per (100   pc³) and/or comparable rates of massive star formation, a nearly continuous string of overlapping regions of ∼10⁴–10⁶  K can be maintained. When lower rates of energy input prevail, intercloud regions will cool and coalesce, forming new clouds. In denser regions, comparable energy input may not be enough to ionize the clouds, except perhaps near the edges of the dense region, for periods as long as 10⁸ years.

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The Outer Heliosphere: The Next Frontiers

E.N. Parker , in COSPAR Colloquia Series, 2001

1 INTRODUCTION

Perhaps the place to begin is the realization that almost all stars evidently have "astrospheres", all too transparent to be seen but undoubtedly as complex as our own heliosphere. The stellar wind is the creator of the astrosphere, just as the solar wind sweeps out the cavity in interstellar space that we call the heliosphere. Thus the origin of the heliosphere and the astrosphere traces back to the hydrodynamics of the million degree solar and stellar coronas. The solar corona appears to be created by the dissipation of mechanical and magnetic energy in the tenuous gas above the dense photosphere. It is that dissipation, evidently in the form of the microflaring in the magnetically "quiet" regions of the Sun, that creates the heliosphere. The staggering complexity of the convective and magnetic machinations on all scales down into the unresolved microstructure of the solar activity gives some idea of the mystery of the stellar corona and astrosphere. Indeed, the mystery does not stop with the microflaring, for we are in the dark as to the origin of the fibril magnetic fields that seem to drive the system from below the visible surface. With the variety of stellar types and circumstances that may be presumed to create stellar winds and astrospheres, the inquiry into the heliosphere and the extrapolation to other stars is bewildering.

The first primitive model of the heliosphere was sketched some 45 years ago, and the subject has come a long way since that time with the advent of the space age. We begin, then, by noting that the heliosphere evidently has been in place since the formation of the Sun and Earth some 4.6 × 10⁹ years ago. Unknown to classical astronomy, the heliosphere remained "silent" until the advance of technology and science first began to uncover its effects. Only in the last half century have we appreciated its existence. Then once we ventured into space the "silent" heliosphere became noisy indeed. There are a number of terrestrial effects, but in the early years they were more puzzling than informative. Some effects are obvious, e.g. the aurora, while others e.g. geomagnetic fluctuations, cosmic ray variations, etc. are detected only by scientific instruments. It was the geomagnetic storm that a century ago first suggested bursts of "solar corpuscular radiation" from the Sun, consisting mainly of protons and an equal number of electrons to provide electrical neutrality. Otherwise space was regarded as a hard vacuum capable of supporting unlimited electric potential differences, at the same time that the zodiacal light was interpreted as sunlight scattered from about 500 free electrons/cm³ at the distance of Earth (1   AU). Then about half a century ago Biermann's ([1], [2]) studies of the anti-solar acceleration of comet tails led to his fundamental pronouncement of the perpetual universal emission of solar corpuscular radiation. The velocity of the solar corpuscular radiation had long been estimated at 10³  km/sec, from the time delay of a couple of days between the flaring on the Sun and the impact of the corpuscular radiation against the outer boundary of the geomagnetic field. Biermann inferred from the measured anti-solar acceleration of gaseous comet tails that the number density of the solar corpuscular radiation at the orbit of Earth is in excess of 103 electrons and ions per cm 3, later revised downward to perhaps as little as 500/cm a based on resonant charge exchange with the cometary atoms. This density seemed to be confirmed by the comparable interplanetary electron density inferred from the intensity of the zodiacal light, considered at that time to be Thomson scattering of sunlight by free electrons. So the solar corpuscular radiation was powerful stuff. Its impact against the geomagnetic dipole field was calculated to confine the field to a distance of about five Earth's radii on the sunward side.

Leverett Davis ([6]) conceived the first sketches of the heliosphere, reproduced in Fig. 1, based on Biermann's declaration of universal solar corpuscular radiation. Davis referred to it as the "cavity in the galactic magnetic field", the term heliosphere originating only thirteen years later in an article by A. J. Dessler. From the existing estimates of the density and velocity of the solar corpuscular radiation Davis suggested that the corpuscular radiation pushed back the interstellar gas and field to a radius of the order of 200   AU. He recognized that the radius of the heliosphere would vary with the 11-year magnetic cycle of the Sun, and he suggested that the varying size of the heliosphere was responsible for the observed variation of the cosmic ray intensity within the heliosphere.

Figure 1. Two sketches of the cavity in the galactic magnetic field (from Davis [6]) with different suggested solar magnetic field forms.

It should be noted here that the origin of the solar corpuscular radiation at the Sun was a mystery at that time, with vague ideas about acceleration in or around the magnetic fields of active regions, sunspots, and flares. Thus the origin was made even more mysterious by Biermann's basic point that the Sun emitted corpuscular radiation in all directions at all times, regardless of the presence or absence of magnetic active regions.

Now by 1956 John Simpson ([25], [13]) had succeeded in determining the energy spectrum of the variation of the cosmic ray intensity with the varying level of activity of the Sun. The variations were first detected by Scott Forbush, using ion chambers, which are sensitive to the muons produced in the atmosphere by cosmic ray protons with energies of 10-20 Gev and up. Simpson invented the cosmic ray neutron monitor which responds to the nucleonic component in the atmosphere, thereby registering the effect of cosmic ray protons down to about 1 Gev, where the time variations are much larger. Using five neutron monitors distributed from the geomagnetic equator to Chicago (at 55° geomagnetic latitude) he exploited the geomagnetic field of Earth as a magnetic spectrometer. He showed that the variations had an energy spectrum that could not be a consequence of an electrostatic potential difference in space, which would be presumed to decrease the energy of each particle by the same amount. Instead, the variations, apart from the bursts of solar cosmic rays from the occasional large flare, showed simply a removal of particles that increased with declining cosmic ray particle energy. He noted that the variations suggested time varying magnetic fields in space.

The great cosmic ray flare of 23 February 1956 showed direct passage of the solar cosmic rays from their origin on the Sun to Earth, arriving promptly at Earth from the direction of the Sun ([12]). Thereafter the solar cosmic ray intensity was observed to decline slowly as if escaping by diffusing through a magnetic barrier beginning at about the orbit of Mars and extending outward to the orbit of Jupiter. The simplest model suggested by the observations was a radial magnetic field extending from the Sun out to the orbit of Mars, with a disordered nonradial magnetic field beyond.

Collectively this indicated a dynamical state of the solar corpuscular radiation and magnetic field in interplanetary space. The challenge, then, was to understand how the corpuscular radiation and interplanetary magnetic field were created by the Sun.

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The Outer Heliosphere: The Next Frontiers

H.-R. Müller , ... B.E. Wood , in COSPAR Colloquia Series, 2001

3 ABSORPTION RESULTS

In Figure 2, we present a calculation of the Ly-α absorption that includes the ISM and the modeled astrospheric component. The histograms are the blue-side wings of high-resolution HST/GHRS spectra of the Ly-α line analyzed by [7]. The absorption feature centered at 1215.37   Å (for λ AND) is due to interstellar deuterium. The dotted line shows the ISM absorption along the line of sight, as determined by [7]. Clearly, there is excess absorption in the spectrum between 1215.4   Å and 1215.5 Å where absorption by the ISM alone is less than observed by HST. In principle, the modeling outlined above allows us to constrain the stellar wind parameters of λ And and ε Ind by varying the stellar wind plasma density, velocity, and temperature, and identifying which model best fits the HST spectra of Fig. 2. Here, we only vary the density (mass loss rate). The dashed and dot-dashed lines in Fig. 2a represent the line profile after both ISM and astrospheric absorption from three λ And models with mass loss rates 5, 10, and 40 times that of the Sun have been included. The calculation of the line profile with both ISM and $40{\dot{M}}_{\odot }$ λ And astrospheric absorption included (triple dot-dashed) results in far too much absorption, especially near 1215.45   Å. The increased opacity in this model compared to the $5{\dot{M}}_{\odot }$ model lies in the greatly increased extent of the hydrogen wall between the AP and BS while densities remain comparable. The neutral temperature in the $40{\dot{M}}_{\odot }$ λ And wall is also higher. The λ And mass loss rate model of $\dot{M}\mathit{=}5{\dot{M}}_{\odot }$ fits the data very well.

Figure 2. Blue side of the HST spectrum of ε Ind (a, left-hand panel) and λ And (b, right-hand panel) as histograms, together with modeled Ly-α profiles at Earth after accounting for ISM absorption alone (dotted), and ISM   +   astrospheric absorption derived from the models discussed in the text.

From [10].

For the ε Ind absorption (Fig. 2b), we employ three models with mass loss rates 0.5, 0.8, and 1. Two profiles "bracket" the data on the high and on the low side between 1215.4   Å and 1215.5   Å leading us to the conclusion that the stellar wind of ε IND has a mass loss rate between 0.5 and 0.8 solar mass loss rates.

The results underscore the importance of the astrospheric neutral H distribution for Ly-α absorption calculations. They also strongly suggest that our model has captured the basic distribution of neutral densities and velocities as well as the basic, effective neutral temperatures that arise at the investigated stars due to the interaction of the LISM with its stellar wind. The reasonably good fit, in turn, supports the principal results, namely the above described morphology of the astrosphere and the estimate of sensible stellar wind parameters. This is the currently most promising method, albeit indirect, to investigate the morphological structure and characteristics of remote astrospheres [10,11].

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X-Ray Astronomy

M.F. Corcoran , in Encyclopedia of Physical Science and Technology (Third Edition), 2003

V.A.1 Normal Stars

Nearly all types of normal (i.e., noncollapsed) single stars generate X-ray emission at some level. Stellar X-ray emission is thought to be dominated by thermal emission arising from one or more of the following processes: magneto-hydrodynamic heating in the outer stellar atmosphere; frictional heating; and shock heating produced by instabilities in an outwardly flowing stellar wind. The sun is the brightest cosmic X-ray source, yet has a rather modest X-ray luminosity compared to other observed stars. The overall flux of X-rays produced by the sun is also small in comparison to the solar emission at other wavelengths. The solar X-ray flux at earth is about 1  ×   10⁻⁸ W m⁻² (in the wavelength range 1–8   ×   10⁻⁸  cm), so that the solar X-ray luminosity is roughly 3   ×   10²⁶  x erg s⁻¹. This is only one ten-millionth of the solar visible-band luminosity. The X-ray emission can vary by factors of 10,000 on short (minutes–hours) time scales. In the sun X-ray emission is produced in the corona, the outermost layer of the solar atmosphere. Figure 4 shows an image of the soft X-ray emission from the solar corona taken by the Yohkoh satellite. The temperature of the corona is about 2   million degrees and varies spatially by about a factor of 3. The corona is heated by the interaction of the stellar magnetic field and the ambient gas. The magnetic field is produced by the so-called dynamo mechanism, in which the rotation of the convective outer stellar envelope generates subsurface electric currents that produce a magnetic field whose axis is roughly aligned with the rotational axis. The magnetic field threads the solar surface and the solar atmosphere. Heating of the corona is due to twisting of the field lines, which locally intensifies the field; magnetic pressure is converted to kinetic gas energy via disruption and reconnection of the magnetic field lines. Differential rotation in the outer solar envelope causes the strength of the magnetic field to vary with an 11-year cycle; this produces a cyclical variation in the numbers of sunspots on the solar surface, and a cyclical variation in solar X-ray emission. At times of sunspot maximum, the solar X-ray flux can increase by a factor of 100, and during solar flares, the solar X-ray emission can increase by a factor of 10,000 for brief periods. Detailed studies of the solar X-ray flux show that the spectrum is dominated by emission lines, indicating the predominance of thermal emission processes.

FIGURE 4. Image of the soft (0.25–4.0   keV) X-ray emission from the sun, from the the soft X-ray telescope on board the Japanese–U.S.–U.K. Yohkoh solar satellite observatory. X-ray emission from the sun is confined to the corona, the outermost layer of the sun's atmosphere. Thermal X-ray emission is produced by conversion of magnetic energy to heat, causing the outer solar atmosphere to reach temperatures of about 2   million degrees. [Credit: ISAS and the Yohkoh SXRT team.]

Solar-type stars (stars that have masses below about five times the sun's mass) are also known X-ray sources, though the nature of the emission in these stars is less well studied than in the solar case. Since these stars have convective outer envelopes and radiative cores like the sun, the emission mechanism is thought to be dominated by magnetic heating in the stellar coronae in a manner similar to the solar corona. Generally the spectrum of the emission has not been well studied, though modest resolution X-ray spectra show the presence of line emission, again suggesting the predominance of thermal emission processes. There is strong evidence that the X-ray luminosity is related to the stellar rotational velocity, indirect evidence of the importance of the stellar dynamo in generating X-rays in these stars. No clear periodicity or cyclical variability in the X-ray flux from solar-type stars has yet been identified, so it is unclear whether the type of activity cycle represented by the sunspot cycle is a general property of all solar-type stars. X-ray emission from solar-type stars is known to be variable, and rapid brief increases in the X-ray emission have often been seen; presumably these "flares" are similar to the better studied flaring activity seen in the sun.

Very massive stars (stars more massive than about 10 solar masses) have radiative envelopes and convective cores, which means that the stellar dynamo is not very effective in these stars. Naively, one expects that these stars would be weak X-ray sources, so it was somewhat of a surprise when the Einstein observatory conclusively detected X-ray emission from a sample of massive stars. It is thought that the emission from these stars is produced, not via magnetic heating, but via shock heating of a portion of the outwardly moving, unstable stellar atmosphere. Very massive stars produce a large flux of radiation at ultraviolet wavelengths and longer; matter at the stellar surface absorbs photon momentum from the UV radiation field and is driven outward as a massive stellar wind. The radiative driving mechanism is very unstable to velocity perturbations, and it is thought that such seed perturbations can steepen into strong shocks, converting wind kinetic energy to thermal gas energy; the relative velocities are such that temperatures should reach millions of degrees and produce observable X-rays. As yet, however, there is no predictive, quantitative physical model to describe the observed emission. Observationally, there seems to be a strong correlation of X-ray luminosity to the total amount of power radiated over all wavelengths, L _x /L _total   ∼   10⁻⁷. Since the total luminosities of these stars are generally hundreds of thousands to millions of times larger than that of the sun, massive stars generate about the same power in X-rays that the sun generates over all wavelength bands. Typically X-ray emission from massive stars shows little intrinsic absorption from the dense stellar winds, suggesting that at least some of the emitting region exists far from the photosphere. Generally the X-ray emission is not strongly variable, again suggesting that the emission is distributed over a relatively large spatial region.

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The Outer Heliosphere: The Next Frontiers

M.A. Lee , H. Fichtner , in COSPAR Colloquia Series, 2001

1 INTRODUCTION

There are two populations of cosmic rays in the outer heliosphere: galactic cosmic rays and anomalous cosmic rays. Both interact strongly with the solar wind. The two populations are shown in Figure 1 which is a schematic diagram of the heliosphere including the Sun, a radially outflowing solar wind with embedded magnetic irregularities, the solar wind termination shock where the solar wind is deflected toward the heliospheric tail (solid circle), the heliopause separating solar and interstellar plasma (inner dashed line), and a possible bow shock where the interstellar plasma flow is decelerated to flow around the heliopause (outer dashed line).

Figure 1. A schematic diagram of the heliosphere showing the solar wind termination shock (solid line), heliopause (inner dashed line), a possible interstellar bow shock (outer dashed line), anomalous cosmic rays (small dots extending out to the heliopause), high-energy galactic cosmic rays (large dots), and low-energy galactic cosmic rays (small dots outside the heliopause).

The galactic cosmic rays (GCR) are denoted by small dots mostly outside the heliopause (low-energy GCRs with energies ≲   300   MeV/nucleon) and large dots throughout space (high-energy GCRs with energies ≲   300   MeV/nucleon). Outside the heliopause the intensity of low-energy GCRs should dominate that of higher-energy GCRs since the GCR differential intensity in the several-GeV energy range is proportional to E^−2.7 . The GCRs originate from outside the heliosphere at supernova shock waves throughout the Galaxy, and also possibly at stellar wind termination shocks, pulsars, or other more exotic objects. They form a sea of energetic particles throughout the Galaxy with a pressure comparable to that of the interstellar plasma and magnetic field. They are partially expelled from the heliosphere by the "sweeping" action of the solar wind in the process known as solar modulation. As shown in Figure 1, the low energy GCRs barely penetrate the heliosheath, that region of space between the termination shock and the heliopause, due to modulation effects. The high energy GCRs penetrate all the way to Earth's orbit and bombard the upper atmosphere, creating the "showers" of nuclear interactions and secondary particles which led to the discovery of cosmic rays and the existence of a modulating environment about the Sun now known as the heliosphere.

The so-called anomalous cosmic rays (ACR) are denoted by small dots mostly within the heliopause in Figure 1. The ACRs should perhaps better be called heliospheric cosmic rays since they owe their existence to the heliosphere. The ACRs originate at the solar wind termination shock in the process of diffusive shock acceleration. They are swept into the heliosheath by the subsonic solar wind flow and may leak across the heliospheric magnetic field at the heliopause into interstellar space. Due to their lower energies of 300 MeV/ion, in spite of their being mostly singly charged ions, they are strongly modulated but do reach Earth orbit during periods of minimum solar activity. Prior to acceleration they arrive in the heliosphere as interstellar neutral gas which flows into the inner heliosphere (shown in Figure 1 as dashed horizontal lines) and is ionized to form interstellar pickup ions, which are advected with the solar wind to the termination shock. At the termination shock the pickup ions, with their suprathermal halo distribution in velocity space, are favored in their injection into the process of shock acceleration.

This paper serves as an Introduction to the Sessions on Galactic and Anomalous Cosmic Rays: Messengers from Outside the Heliosphere.

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Interstellar Molecules, Spectroscopy of*

A.G.G.M. Tielens , in Encyclopedia of Spectroscopy and Spectrometry (Second Edition), 1999

Introduction

The origin and evolution of life has long fascinated mankind. It is directly tied to the abiotic formation and chemical history of the biogenic elements, in particular H, C, N, O, P and S. The evolution of most elements starts with nucleosynthesis inside the fiery cauldrons of stellar interiors. Eventually, these elements are injected into the interstellar medium, either through violent supernova explosions or through more gentle stellar winds, which terminate the life of most stars. Thus, stellar formation, evolution and death are intimately connected, as the ashes of one generation of stars become the building blocks of the next generation. Much of the expelled material is in the form of molecules, ranging from the simple diatomics such as molecular hydrogen (H ₂) and carbon monoxide (CO) to more complex species such as polycyclic aromatic hydrocarbons (PAHs) and fullerenes (C₆₀), or in the form of small (≈100   nm) dust grains – carbon based (i.e. soot and silicon carbide) and silicates (aluminium oxide, olivine, enstatite). In evolution from their birth sites, through the interstellar medium, until their incorporation into new planetary systems, these elements undergo a complex chemical history. Various processes cycle the material from the gas phase to the solid phase and from one chemical compound to another. Understanding this history and its relation to the origin of life is one of the key questions of modern astrophysics.

The most important tool for astronomers to understand the origin and evolution of the molecular universe is spectroscopy. This includes studies of the rotational emission spectrum of the cold interstellar gases in the microwave wavelength region. Near bright stars, vibrational modes can be excited through fluorescence. Electronic spectra are predominantly measured through absorption spectroscopy against a bright background source. Each of these wavelength regions has its own techniques and provides unique information on interstellar molecules. Together they point to a chemically diverse and active molecular universe. Each of these techniques and their results will be briefly summarized.

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Binary Stars

Steven N. Shore , in Encyclopedia of Physical Science and Technology (Third Edition), 2003

II.C Distortions in Photometric and Velocity Curves

Several effects have been noted that distort the light curve and can be used to determine more physical information about the constituent stars.

II.C.1 Reflection Effect: External Illumination

Light from one component of a close binary can be scattered from the photosphere and outer atmosphere of the other, producing a single sinusoidal variation in the system brightness outside of eclipse. This reflection effect is useful in checking properties of the atmospheres of the stars. If the illuminating star is significantly higher temperature, it can also produce a local heating, which alters the atmospheric conditions of the illuminated star. Such an effect is especially well seen in X-ray sources, particularly HZ Her   =   Her X-1, which varies from an F-star to an A-star spectrum as one moves from the unilluminated side to the substellar point facing the X-ray source.

II.C.2 Photospheric Nonuniformities: Starspots

A similar phenomenon has been noted in the RS CVn stars, where it is caused by the presence of large-scale, active magnetic regions, called starsports, on the stellar surfaces. Unlike reflection, these dark regions migrate with time through the light curve as the active regions move with the differential rotation of the stellar envelope, analogously to the motion of sunspots. Chemically peculiar magnetic stars also show departures from simple eclipse profiles, because of locally cooler photospheric patches, but these appear to be stable in placement on the stellar surface.

II.C.3 Circumstellar Material

The presence of disks or other circumstellar matter also distorts the light curves and can alter the radial velocity variations as well. In Algol systems, this is especially important. The timing of eclipses indicates a circular orbit, while the radial velocity variations are more like that of a highly eccentric one. The explanation lies in the fact that here is considerable optical depth in the matter in the orbit, which results in the atomic lines producing a distortion in the radial velocity variations. Many of the W Serpentis stars show this effect. It is most noticeable in eclipsing systems because these present the largest path length through material in the orbital plane. In some cases, atmospheric eclipses can also distort the lines because of stellar winds and convection cells intercepted by the line of sight. These motions, however, are generally small compared with the radial velocity and so alter the photometry (light-curve instabilities during eclipse are well marked in the ζ Aur stars) but do not seriously affect the radial velocity determinations.

II.C.4 Ellipsoidal Distortions: Tidal Interaction

If the stars are close enough together, their mutual gravitational influences raise tides in the envelopes, distorting the photospheres and producing a double sinusoidal continuous light variation outside of eclipse. Many of these systems also suffer from reflection-effect distortion, so there are many equivalent periods in these systems, depending on whether or not they eclipse.

Departures from symmetric minima should accompany expansion of the stars within their tidal surfaces. As the photosphere comes closer to the tidal-limiting radius, the Roche limit, the star becomes progressively more distorted from a symmetric ellipsoid and the photometric variations become more like sinc curves. An additional feature is that as the stars become larger relative to the Roche limit they subtend a greater solid angle and display increasing reflection effect from the companion.

II.C.5 Limb Darkening

Stellar surfaces are not solid, and they have a continuous variation in surface brightness as one nears the limb. This effect, called limb darkening, is produced by the temperature gradient of the outer stellar atmosphere compared with the photospheres. The effect of limb darkening on light curves is to produce a departure from the behavior of simple, uniform spheres most notable in the softening of the points of contact during eclipse. It is one of the best ways available for measuring the temperature gradients of stellar atmospheres.

II.C.6 Apsidal Motion: Orbital Precession

The additional effect of the tidal distortion is that the stars are no longer simple point sources, but produce a perturbation on the mutual gravitational attraction. The departure of the gravitational potential from that of two point masses produces apsidal motion, the slow precession of the line connecting the two stars. This rate depends on the degree of distortion of the two stars, which in turn provides a measure of the degree of central concentration of the stars. Such information is an important input for stellar evolutionary models. One system that has been especially well studied is α Virginis (Spica). An additional source of apsidal motion is the emission of angular momentum from the system, and the presence of a third body.

II.C.7 Third Light

Either because of circumstellar material in the orbital plane, which is not eclipsed but which scatters light from the binary components, or because of the presence of a faint third body in the system that is unresolved, some additional light may be present at a constant level in the eclipsing binary light curve. Frequently, high-resolution spectroscopy is able to reveal the lines of the companion star, as in Algol, but often it remains a problem to figure out the source for the nonphotospheric contributions to the light curve. This is simply added as an offset in the determinations of eclipse properties in most methods of light-curve analysis. Such emission may also arise from shocks in accretion disks and from intrinsic disk self-luminosity.

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The Outer Heliosphere: The Next Frontiers

Nikolai V. Pogorelov , in COSPAR Colloquia Series, 2001

2.1 Current achievements

The presence of the solar and interstellar magnetic fields necessitates solving the MHD equations to model the solar wind (SW) interaction with the local interstellar medium (LISM). The influence of magnetic field is important if the magnetic pressure becomes comparable with the dynamic pressure of the flow. The magnitude and direction of the interstellar magnetic field is uknown. We can only say that a strength of about 1.6 μG is consistent with observations of polarization and pulsar dispersion measurements [15]. Estimates of the heliosphere confinement pressure give an upper limit of 3–5 μG (see [10]). Concerning the direction of the magnetic field in the local interstellar cloud, we can only refer to [16] which states that if the cloud surface is perpendicular to the direction of the gas flow, the direction of the magnetic field is approximately parallel to the cloud surface. These estimates lead us to the conclusion of the importance of MHD modeling.

The first numerical study of the stellar wind interaction with the magnetized LISM was performed in / 17,25/ by a rather viscous flux-splitting method for the case of B _∞ || V _∞, where the subscript ∞ refers to the LISM parameters. The authors, actually, performed parametric calculations for various ratios of dynamic pressures and stagnation temperatures and some of them were very far from the SW–LISM interaction case. In [3] very accurate shock-fitting calculations of the upwind region of the interaction were performed and some of the results obtained in [17] were criticized. As later turned out [28], this mainly concerned the case of an irregular interaction in which parallel fast MHD shocks were nonevolutionary (structurally unstable), while singular (switch-on) shocks could not exist due to the symmetry restrictions.

It was noticed in [45] that the toroidal component of the interplanetary magnetic field (IMF) can be important. Three-dimensional calculations with the presence of IMF for B _∞ || V _∞ were reported in [46].

Systematic shock-capturing MHD calculations of axisymmetric problems for various magnitudes of the interstellar magnetic field (ISMF) were made in /28,31/. Both regular interaction with a single bow shock and irregular interaction with additional discontinuities were discovered. The general configuration of the flow pattern corresponding to the two-shock model [4] is shown in Fig. 1 for the following SW and LISM parameters: n_e   =   7   cm⁻³, V_e   =   450   km   s⁻¹, M_e   =   10, n_∞   =   0.07   cm⁻³, V _∞  =   25   km   s⁻¹, and M _∞  =   2. Here the subscript e refers to the SW parameters at Earth's distance from the Sun. M stands for the Mach number. The dimensionless value of the magnetic field is specified via the Alfvén number $A_{\infty }=V_{\infty }/\sqrt{B_{\infty }^2/4\pi {\rho }_{\infty }}$ . In the presented case $A_{\infty }=\sqrt{2}$ (see [28]). The abbreviations BS, TS, and HP correspond to the bow shock, the heliospheric termination shock, and the heliopause, respectively.

Figure 1. General configuration of the heliosphere: density (below the symmetry axis) and total pressure logarithm isolines [28]. The LISM flow is directed from right to left. The Sun is at the origin.

The influence of the ISMF direction on the shape of the global heliopause was studied in /27],[36/. The results turned out to be consistent with the MHD modeling of the three-dimensional heliopause on the basis of the Newtonian approximation [14]. The distribution of the streamlines and magnetic field lines in the symmetry plane for the cases with A _∞  =   2 for two different angles between ISMF and LISM velocity are shown in Fig. 2 (see [27]). All previously cited calculations disregarded the charge exchange processes occurring among neutral and charged particles which are extremely important in the SW–LISM interaction. The IMF was also neglected.

Figure 2. Streamlines and magnetic field lines for A _∞  =   2 in the symmetry plane. The angle between B _∞ and V _∞ is equal to 90° (left) and 45° (right).

Three-dimensional calculations for various direction of ISMF, with IMF in the form of the Parker nominal spiral and with the simplified treatment of the charge-exchange processes were performed in [23]. McNutt at al. [26] investigated three-dimensional, both purely hydrodynamic and MHD flows, taking into account the latitudinal variation of the SW velocity and density in accordance with the Ulysses data.

The solar cycle dependence of the heliospheric shape was studied in [42] on the basis of a global MHD simulation of the time-dependent SW interaction with the magnetized LISM. Charge exchange processes were, however, neglected. This is not very reasonable, since the one-dimensional results /37/ and /44/ show that even if magnetic field is added, pick-up ions still remain a dominant factor of the interaction. This is in agreement with the recent results [1] where self-consistent axisymmetric MHD calculation were performed with Monte Carlo modeling of the neutral particle motion.

The accuracy of numerical calculations can be substantially increased if we reduce the size of the computational region. Otherwise, the resolution of discontinuities will be rather poor [36]. To do this, we must shift the boundary conditions from infinity to the finite distance from the Sun and state nonreflecting boundary conditions at it. Below we consider an efficient approximate procedure created for this purpose.

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Source: https://www.sciencedirect.com/topics/physics-and-astronomy/stellar-wind