from th e UV massloss rates of G arm any et al. (1981) and G arm any and Conti (1984) and the radio massloss rates of A bb ott (1985), and calculated the integrated H a lum inosity using a CL velocity law w ith various values of (3. For each value of the velocity law index, Drew was thus able to derive an equation relating M to L ( H a ), and the stellar photospheric radius. These formulae were then used to calculate massloss rates for seven stars w ith reliable radio M determ inations by using the values of L ( H a ) given by Leitherer (1988). It was found th a t the best agreement between the two m ethods was given when = 1.5 rath er than the theoretically preferred value of ~ 0.7 which, on average, gave values of M a factor of 2 higher th an th e radio determ inations. Since Leitherer’s (1988) massloss rates agreed w ith those determ ined by other m ethods Drew compared the theoretical 0 star H a massloss rates, calculated for (3 = 0.7, w ith those determ ined by Leitherer and found th a t the factor of two discrepancy was still present even when the greater num ber of less luminous stars in Leitherer’s sample were included. M ost observational and theoretical studies indicate th a t 0 .7 < /? < 1 .0 (e.g. Friend and A b b o tt, 1986; Pauldrach et al., 1990) so Drew interpreted her results as possible evidence for th e presence of instabilities in the stellar winds. Drew calculated th a t in a sm ooth outflow the H a emissivity would peak inside ~ 1.3i£*, whereas th e initial modelling of Owocki et al. (1988) has suggested th a t wind instabilities will cause compressions in the flow which have greatest density between ~ 1 .2-2.3 R*. Drew argued th a t these density enhancem ents would increase the H a emissivity by factors of a few and could produce a sufficient increase in L ( H a ) to explain the over-large values of M found from the H a observations. In view of this, Drew gave th e following expression for determ ining massloss rates from L ( H a ) which was calculated for (3 — 0.7 and has been qualitatively corrected for th e effect of density inhomogeneities
photom etry and I R A S LRS spectra, together with I R A S broad band photom etry at 12, 25, 60 and 100/nn. Stars chosen are those w ith inform ation on their velocities and distances. Some of them are optically thin, i.e., the 9.7/zm silicate feature is in emission, while some are optically thick, w ith the silicate feature in absorption. Some stars were chosen because they have evidence of ice in their outflows. Models for com bination grains are successfully obtained for the first time. Dust massloss rates, one of the results from these models presented in chapter 3, are outflow rates relatively close to central stars. The same is tru e for massloss rates derived from ratios of I R A S fluxes a t 25 and 12/nn, while 60/im photom etry can be used to obtained dust massloss rates farther out in the shells. It is shown th a t estim ated dust massloss rates using various m ethods are in reasonable agreements. Also, gas massloss rates are calculated using expressions from the literatu re, and are compared to dust massloss rates in order to estim ate d u st-to -g a s mass ratios for these stars. Radio observations of OH masers can be used to obtain gas massloss rates (e.g., Bowers et al. 1983) out to a few 1016cm, and the sub-m illim etre CO J = l - 0 and J = 2 - l lines have also been shown to be a very useful tool in estim ating gas massloss rates in the outerm ost region of circumstellar shells (e.g., K napp & Morris 1985). D ust massloss rates estim ated from chapter 3 reflect the more recent period of massloss from stars than massloss rates derived from radio and sub-m illim etre observations, since infrared emission arises closer to the central stars ( ~ 1015cm). Results show th a t silicate dust models used are reasonable, and th a t the main uncertainty in estim ating dust massloss rates comes from the estim ated distances to the stars. Comparisons between dust and gas massloss rates lead to conclusions th a t either the d u st-to -g a s mass ratio is not a constant from sta r to star, or th a t massloss rates vary w ith time. If massloss rates do vary w ith tim e, a radiative transfer code made to cope w ith this would be valuable.
The time-averaged radio spectrum of Cyg 0B2 No.9 displays a very flat spectral index, of Û! = 0,06 ±0.03. At specific epochs, however, the star has exhibited negative, positive, or even complex spectral indices. This might be explicable in terms of (variable) nonthermal emission, with an underlying thermal contribution from the stellar wind (Cyg 0B 2 No.9 is an extremely luminous star, so an observable thermal contribution would be expected). At two epochs. May 9th, 1983, and May 1st, 1993, Cyg 0B2 No.9 exhibited a positive spectral index, of a = 0.56 ± 0.36 and a = 0.59 ± 0.06, respectively. These distributions are indicative of free-free emission, perhaps suggesting that the nonthermal component was absent at these times. Coincident with these events were significant decreases in absolute flux levels (see Table 6.1), again suggestive of an ‘absent’ nonthermal contribution. Interestingly, the level of intensity of the 6 cm flux was a factor ~ 4 greater in May 1993, which, assuming that the observed emission at these epochs was thermal, implies an increase in mass-loss rate of a factor of ~ 3, as compared to May 1983. If M is calculated using the lower. May 1983 flux (this is the lowest 6-cm flux recorded for Cyg 0B 2 No.9), the result is still a factor ~ 2 higher than the H a mass-loss rate derived in §4 (see the notes to Table 6.3, on page 176, for a comparison of the values). It would be desirable to obtain multiwavelength observations (including Ha) of Cyg 0B2 No.9 at other epochs of decreased flux (February 9th, 1982, and August 22nd, 1983, also saw reduced emission), to investigate whether or not the nonthermal contribution ever disappears entirely.
Figure 2 compares the empirical mass distribution (histogram) based on the extended sample of 22 stars with the expected range of possible sdB masses from various theories for the formation of such stars. The red boundaries correspond to the calculations for single star evolution, in which a sdB star is formed from a red giant star experiencing a strong massloss on the RGB . These authors predict a narrow range of possible masses, between around 0.41 M and 0.52 M . The upper limit is rather sharp and corresponds, in the computations of , to the mass limit above which a star will ascend the AGB instead of becoming an EHB star. The lower limit is more fuzzy and corresponds to the still uncertain value of the mass required to ignite helium in the stellar core (such a minimum mass depends, e.g., on the metallicity of the progenitor). The blue curve on Fig. 2 is the weighted mass distribution for the 3 binary evolution scenarios of [5, 6], including observational selection effects. In this curve (coming from Fig. 22 of ), almost all sdBs with masses above 0.5 M come from the two He-WD merger channel.
The second set of ZDI-driven MHD models is presented here, which includes the simulated stellar winds and inner astrospheres of the systems of interest. As in the previous chapter, solar simulations are compared with spacecraft data and used to evaluate the reliability of the simulations. A fundamental result of this chapter is the characterisation of the structure of the stellar wind and its physical properties, in connection with the underlying distribution of the surface magnetic field. Moreover, a prediction for the massloss and angular momentum loss rates of these systems is obtained. These results are compared and discussed in the context of previous observational and numerical works. Additional simulations are carried out for one of the systems (HD 1237), to investigate in more detail the conditions experienced by the exoplanet at various locations of its orbit. This is done by assuming a magnetosphere around the planet, which interacts self-consistently with the incident stellar wind. This analysis reveals a dominant role played by the stellar wind density (over the velocity), in the process of particle injection through the shield provided by the exoplanet magnetosphere. Following previous studies, the amount of magnetospheric radio emission from the planet in this system is refined, and the possibilities of its detection using current and future instrumentation are discussed. Finally, the stellar wind properties at the inner edges of the Habitable Zones (HZ) of these stars are calculated, which can be used in further studies of astrobiology and climate modelling of exoplanetary systems.
Let us now consider how rotation changes the picture. Fig. 3, right-hand panel, shows the evolutionary tracks of the Z = 0.014 rotating tracks in a similar way as in the left-hand panel. The changes brought by rotation are modest. This is expected because of two facts: first, in this high mass range, the evolution is more impacted by massloss than by rotation, secondly, stars are already well mixed by the large convective cores. One notes however a few differences between the non-rotating and rotating models. One of the most striking is the fact that the models during their O-type phase evolve nearly vertically when rotation is accounted for. This is the effect of rotational mixing which keep the star more homogeneous than in the non-rotating cases (although, as underlined above, already in models with no rotation, due to the importance of the convective core, stars are never very far from chemical homogeneity). As was the case in the non-rotating tracks, the O-type star phase corresponds to an upward displacement when time goes on in the HRD for the 150 M model, while, it corresponds to a downwards displacement for the three more massive models. One notes finally that lower luminosities are reached by the rotating models at the end of their evolution (decrease by about 0.3 dex in luminosity, thus by a factor
The use of IR for deriving stellar and wind parameters has become a necessity, as many stars are obscured. Many diagnostics are available in this range, however, so that using IR does not necessarily imply restrictions of the output quality. Line EWs and the appearance of the spectra can be used to derive temperature or spectral types, though morphological spectral types so inferred are less precise. This range is also helpful for mass-loss diagnostics, especially for low-massloss rates. Two caveats, however, must be noted: the NIR lines form in the wind acceleration zone and they are very sensitive to NLTE and 3D effects - they are thus extremely sensitive to modelling details. For most of the cases, observing K-band is sufficient, but diagnostics improve with J and H spectra, or - even better - if IR is complemented by optical and UV data (terminal velocities, are not constrained by IR, for example). Forbidden lines observed in the IR may arise
We consider a viscous flow enclosed within a rotating spherical shell. We suppose the centrifugal acceleration does not deform it. We study the interaction of the com- bined e ff ect of the inertia with rotation in a stably strati- fied environment, leaving aside magnetic fields (, , ), internal gravity waves (, ) and anisotropic turbulence (, , ) as a first step. Because we are interested in the evolution of the fluid on secular timescale, we solve as a first step the steady equation of the vortic- ity combined with the equations of energy and continu- ity using the Boussinesq approximation (, , ). This approximation retains density variations only in the buoyancy term taking into account both the gravity and centrifugal accelerations and implies that the stratification is only taken into account through the Brunt-Väisälä fre- quency profile (positive in a radiative zone). It is given as an input of the simulations using the 1D MESA stellar evolution code () to generate ZAMS low massstars radiative core models (we use a metallicity of Z = 0.02, and a mixing length theory parameter of α MLT = 2) and
2.5. DISCUSSION 53 The metallicity distribution is best represented by either an accretion or a leaky box model when corrections for the sampling bias are taken into account (see Figure 2.7). The only differences between the two are in the low–metallicity tail. Carrera et al. (2008) proposes a model featuring gas moving both in and out of the system as the best representation, with free parameters of α = 1.2 and λ = 0.4 − 0.6. This is equivalent to a net inflow with α = 0.6 − 0.8, conditions not significantly different to the closed box model when representing the corrected metallicity distribution. Interestingly, the age–metallicity relation is not well– reproduced by any simple model, because of the large scatter in metallicity at intermediate ages. Increasing the yield value used in the models (to be closer to the value used by Carrera et al., 2008) better reproduces the peak metallicity in the distribution, but shifts the AMR too far to the young, metal–rich corner. The SFH predicts that only a small number of RGB stars older than 6 Gyr should be observed, in agreement with observations, but the high metallicity of stars from 6–12 Gyr is not well–fit by the simple models. The cohort of metal–poor stars at ages 2–6 Gyr is also problematic for the chemical evolution models. It is likely that both more accurate age estimates for the RGB stars are required and more complex models need to be considered, but this is beyond the scope of this research.
temperature, metallicity and surface gravity (e.g. Santos et al. 2005; Sousa et al. 2006). Currently, the knowledge of the photometry (including bolometric correction) and parallax allows the luminos- ity (L) to be determined and then the mass can be estimated. This methodology is potentially interesting for nearby stars, where the uncertainty in the distance (and thus in the luminosity) is more likely to be small. Errors affecting this method include the usually high uncertainties in the spectroscopic surface gravities, or for the case of stars at more than 50 pc, the errors in the measured parallax (for some stars, this problem will be solved with the Gaia mission). The present accuracy on the mass determinations using this method is not better than 10–20 per cent (Sousa et al. 2011a).
The apsidal rotation (or periastron rotation) of close binary stars is a result of their non-Keplerian movement which originates from the non-spherical form of stars. This non-sphericity has been produced by rotation of stars around their axes or by their mutual tidal effect. The second effect is usually smaller and can be neglected. The first and basic theory of this effect was developed by A.Clairault at the beginning of the XVIII century. Now this effect was measured for approximately 50 double stars. According to Clairault’s theory the velocity of periastron rotation must be approximately 100 times faster if matter is uniformly distributed inside a star. Reversely, it would be absent if all star mass is concentrated in the star center. To reach an agreement between the measurement data and calculations, it is necessary to assume that the density of substance grows in direction to the center of a star and here it runs up to a value which is hundreds times greater than mean density of a star. Just the same mass concentration of the stellar substance is supposed by all standard theories of a star interior. It has been usually considered as a proof of astrophysical models. But it can be considered as a qualitative argument. To obtain a quantitative agreement between theory and measurements, it is necessary to fit parameters of the stellar substance distribution in each case separately.
alone. Using the Güdel – Benz relation and the observed X-ray luminosity of V374 Peg ( L X = 10 28.44 ; Hünsch et al. 1999 ) , V374 Peg ’ s radio luminosity from gyrosynchrotron emission alone should be L n , R = 10 12.94 . Using the distance to V374 Peg ( d = 8.93 pc; van Leeuwen 2007 ) , this luminosity corresponds to a radio ﬂ ux of F X ~ 0.08 mJy. From the VLA observations ( Figure 3 ) the observed radio ﬂ ux is at least one order of magnitude higher than this value, suggesting that gyrosynchro- tron emission is a negligible contribution to the total radio ﬂ ux from V374 Peg. Note that there is uncertainty in the Güdel – Benz relation, particularly for low-massstars and ultracool dwarfs that appear to lie above this relation. Simultaneous VLA and Chandra observations of the Orion Nebula Cluster by Forbrich et al. ( 2017 ) enabled these authors to search for correlations between extreme radio and X-ray variability from young stellar objects. They found 13 radio sources, all of which also exhibited X-ray variability. Multi-epoch radio, optical ( including H α) , UV ( Swift ) , and X-ray ( Chandra ) observations of the UCD binary NLTT 33370 AB by Williams et al. ( 2015 ) found periodic modulation in the radio and optical and plausible modulation in H α and the UV. Comparing simulta- neous X-ray light curves with radio observations may help assess the relative contributions of radio emission through ECM and gyrosynchrotron processes. If the dominant source of radio emissions is through the ECM instability as modeled here, the radio and X-ray light curves should be anti-phased; however, if the dominant emission process is gyrosynchrotron emission then the light curves should be phased.
for higher mass stellar members of USco with spectral types B–M5 (triangles). The red solid line indicates a linear ﬁt to the data, computed with 3σ clipping of outliers and presumed to represent photospheric colors. RMS residuals of the ﬁt are domi- nanted by small uncertainties measured for the stellar photometry; thus, the limiting uncertainties are not ﬁt residuals but the photometric uncertainties of the low mass brown dwarfs. I therefore adopt a threshold to deﬁne an excess source corresponding to 3 × the average ([4.5]-[8.0]) error for the brown dwarfs, or ∼ 24% above photospheric colors. The four brown dwarfs with a [3.6]-[8.0] color excess also have an excess at [4.5]-[8.0] colors, and exhibit similar [4.5]-[8.0] colors to those computed for K- and M-type stars with disks identiﬁed by Carpenter et al. (2006). Based on the adopted criteria listed above to identify infrared excess sources, the disk frequency for M stars (M0–M5) in the study by Carpenter et al. (2006) is 16/101 ( ∼ 16 +5 − 4 %). I measure a disk frequency for brown dwarfs of 15 +12 − 7 %. Using a two-tailed Fisher’s Exact test, these frequencies are indistinguishable from each other implying a near-constant fre- quency in USco of stars with disks compared to stars without disks spanning a mass range from ∼ 0.5 M to ∼ 0.02 M . Thus, the observations I presented here add to the growing evidence that substellar formation and early evolution must proceed through very similar physical processes as those that form and shape low massstars.
an important element in the chemical evolution of the ISM. Met- als resulting from stellar nucleosynthesis are returned to the ISM, either as gas and solid grains condensed during the later stages of stellar evolution; they can later be destroyed and in- corporated into new generations of stars. Elements are injected into the ISM at di ff erent rates (we refer to e.g., Dwek et al. 2009). For instance, atomic bounds of carbon are most likely responsible for the emission features, whose carriers are often identified with Polycyclic Aromatic Hydrocarbons (PAHs); they might originally have been produced in AGB stars, whose typical lifetimes are a few Gyr (e.g., Marigo et al. 2008; Cassarà et al. 2013; Villaume et al. 2015; Simonian & Martini 2017). Other el- ements may be synthesized in more massive stars, which die as supernovae on shorter timescales. Thus, the abundance of the mid-IR (MIR) emission feature carriers and the dust-to- gas mass ratio (DGR) are expected to vary with the age of the stellar populations and correlate with the metal (e.g., oxy- gen and nitrogen) abundance of the gas (e.g., Galliano et al. 2008). However, dust formation results from a long chain of poorly-understood processes, from the formation and injection into the ISM of dust seeds formed in the atmospheres of AGB stars or in SNe ejecta, to grain growth and destruction in the ISM (e.g., Valiante et al. 2009; Asano et al. 2013; Mattsson et al. 2014; Rémy-Ruyer et al. 2014; Bocchio et al. 2016). Based on Draine (2009); only ∼10% of interstellar dust is directly formed in stellar sources, with the remaining ∼ 90% being later con- densed in the ISM. However, Jones & Nuth (2011) found that the destruction e ffi ciencies might have been severely overesti- mated. They concluded that the current estimates of global dust lifetimes could be uncertain by factors large enough to call into question their usefulness (we also refer to Ferrara et al. 2016).
Abstract. Low massstars may lose their envelopes in the first giant branch (RGB) or the asymptotic giant branch (AGB) via envelope ejection (i.e. superwind). The envelope loss of AGB stars leads to the formation of carbon-oxygen (CO) white dwarfs (WDs), while the envelope loss of AGB stars may lead to the formation of helium WDs. We mainly focus here on where a RGB/AGB star loses its envelope during its evolution and we show the inital - final mass relation. We also propose a possible channel for the formation of single hot subdwarf stars, in which an old metal-rich RGB star with positive envelope binding energy may lose its envelope and the naked helium core gets ignited to become a hot subdwarf. We also review the well- established Han et al. scenario for the formation of hot subdwarf stars, in which binary interactions lead to the formation of both single and binary hot subdwarfs. By detailed binary evolution calculations, we show that PG 1018-047, a hot subdwarf binary with a main sequence companion and a very long orbital period of 756 d, is explained naturally from the stable RLOF channel in the Han et al. scenario.
photodissociation is sufficiently large, the atmosphere can be pushed out of thermochemical equilib- rium (see, e.g., Moses, 2014; Miguel & Kaltenegger, 2014; Hu & Seager, 2014). This effect has been observed in the cases of the hot Jupiters HD 209458b, HD 189733b, and WASP-12b (Vidal-Madjar et al., 2003; Linsky et al., 2010; Lecavelier Des Etangs et al., 2010; Fossati et al., 2010a,b, 2013), and the hot Neptune GJ 436b (Kulow et al., 2014; Ehrenreich et al., 2015). It should be noted that, although seemingly damaging, the atmospheric effects caused by high energy radiation might also influence the emergence and evolution of life. Paradoxically, UV radiation can both damage (Voet et al., 1963; Matsunaga et al., 1991; Tevini, 1994; Kerwin & Remmele, 2007) and aid in synthesizing (Senanayake & Idriss, 2006; Barks et al., 2010; Ritson & D Sutherland, 2012; H Patel et al., 2015) many molecules critical to the function of life on Earth. Thus, the effects of stellar irradiation on planetary atmospheres warrants the comprehensive spectroscopic characterization of low-massstars at short wavelengths.
The thermogravimetry of the complex of Mn(II) shows the loss of the ligand in three steps of massloss. In a fourth step it loses part of the bromine content leaving a residue that is part of the bromine content plus the metal content. The complex of Fe(II) shows the loss of the ligand in three steps. Part of the bromine content is lost to- gether with part of the ligand in the third step of massloss. Part of the bromine content is lost in the fourth and fifth steps of massloss leaving a residue that is part of the bromine content plus the metal content. The complex of Co(II) shows the loss of ligand in three steps of massloss. Part of the bromine content is lost together with part of the ligand in the third step. Part of the bromine content is lost in a fourth step of massloss leaving a resi- due that is part of the bromine content plus the metal content. The complex of Ni(II) shows the loss of the ligand in the first step of massloss follow by the loss of part of the bromine content in a second step of massloss leav- ing a residue that is part of the bromine content plus the metal content. The complex of Cu(II) shows the loss of the ligand in the first step of massloss together with part of the bromine content. The rest of the bromine content together with part of the metal content is lost in the second step of massloss leaving a residue that is part of the metal content. The complex of Zn(II) shows the loss of the ligand in three steps of massloss. The bromine con- tent is lost together with part of the metal in the third step of massloss. Part of the metal content is lost in the fourth step of massloss leaving a residue that is part of the metal content. The complex of Cd(II) shows the loss of the ligand in two steps of massloss follow by the loss of the bromine content and part of the metal content in the third step of massloss leaving a residue that is part of the metal content. Figure 2 presents the TG/DTG curve of the Co(II) complex. The DSC curves of the complexes are consistent with the TG data. They present endothermic peaks due to the elimination of part of the ligand or part of the bromine content alone or together with part of the ligand. An exothermic peak is observed in the DSC curve of the Fe(II) complex due to the de- composition of the complex. Figure 3 presents the DSC curve of the Co(II) complex. Table 3 presents the thermoanalytical data for the complexes.
this point, and to set the values on a convenient scale. In Figure 17, we have reproduced a version of Figure 24 of Albrecht et al. ( 2012 ) , showing the relative tidal dissipation timescales for all hot Jupiters with measured spin – orbit angles. We have excluded the same systems they did and added nine recent projected obliquity measurements, including that of Kepler-432b. Their result still holds true: systems with short tidal dissipation timescales are well aligned, while those with long tidal timescales display a wide range of obliquities. We also note that recent work by Valsecchi & Rasio ( 2014 ) , which includes a more detailed treatment of the convection and stellar evolution of each star, similarly concludes that the observed obliquity distribution can be explained by tidal evolution. While the angle measured for Kepler-432b is the line-of-sight projection rather than the sky-plane projection, and the ef ﬁ ciency of realignment may be different for evolved stars given their different internal structure, Kepler-432b does have a short tidal timescale, and it sits in a region with hot Jupiter systems that are mostly well aligned. We interpret this result as evidence that the spin axis of Kepler-432 has been realigned to the orbit of the inner planet by the same mechanism responsible for hot Jupiter realignment.
According to most authors [e.g., 29], the main driver of blue HB morphologies, in the face of abundance variations in (mainly) O, Na, Mg, and Al, is the associated level of He enhancement that is brought about in the course of the operation of the Ne-Na and Mg-Al proton-capture cycles, which give rise, for instance, to the well-known O-Na anticorrelation. Uncertainties in the relevant nuclear reaction rates can be quite large, reaching orders of magnitude in some cases [e.g., 49–51]. In fact, the precise level of He enhancement that may be associated with a given degree of O depletion/Na (and Al) enhancement is not known a priori. As an example, in Figure 5 (left panel) the predicted correlation between [Na/Fe] and [O/Fe] [adapted from 57] is shown for two well-known GCs, namely M13 (NGC 6205) and NGC 2808, both of which harbor large numbers of extreme HB stars. The right panel in the same figure shows the associated trend between the helium abundance Y and [O/Fe]. It is especially noteworthy that extremely low levels of [O/Fe] may be reached, with only a minor level of associated He enrichment. In addition, closely the same [O/Fe] and [Na/Fe] ratios may imply widely different He enrichment levels. Last but not least, note also that some of the variation in the observed abundance ratios may be evolutionary in nature (i.e., reflecting mixing processes operating in the interiors of the red giants), as recently demonstrated by . It is important to keep these points in mind, before arriving at conclusions regarding the level of He enrichment that may be present in a given population, based solely on measurements of light-element abundance ratios.
He-4 fusion heats the core rapidly and a ‘helium flash’ takes place, causing the core to expand. This lowers core temperature and reduces the total energy output; then the outer layers contract and the star’s temperature increases a bit. After about 100 million years, the star fuses all its core helium into C-12. Then a helium fusion shell forms around the core, and hydrogen fusion shell remains around that. Then the star again becomes a red giant and remains like that for a few million years with its outer layers continuing to expand. Well-known examples of red giants are Aldebaran in the constellation Taurus, Arcturus in the constellation Boötes and Gamma Crucis in the constellation Crux. These are some of the brightest stars in the night sky.