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83113 (2000)) | $83443$ \cite{mayor83443}) \ref{table:parameter}. |
1.. ο. dliucjifdt=l/r1963.. 1981)). dq22:«107e. hoOX 1995)). (Q—3«107 1966)). 7~23«105 6.5 esic2s. 83113e 831130's 2000)). (2000})) /. £. ei £ J 109615 al.20003). 281165 1981)). 19958)). 3)). ἐν | $e_1 = 0.079 \pm
0.008$ $e_1$ $d\ln e_1/dt = -
1/\tau$\cite{goldreich}, \cite{hut}) $R_1\approx 7\times
10^4\km$$m_1\approx 2\times 10^{30}\g$ $k_{2} \approx 0.5$ \cite{yoder}) $Q
\sim 3 \times 10^5$ \cite{goldreichsoter}) $\tau \sim 3\times 10^8$ $6.5\,$ $v \sin i < 2 \km/\s$ $83443c$ $83443b$ \cite{mayor83443}) \cite{mayor83443}) $J$ $E$ $e$ $E$ $J$ $109648$ \cite{jha}) $284163$ \cite{griffin}) \cite{kiseleva}) \ref{sec:linear}) $i$ |
shell, increases by a factor ~ 40 in ~ 35 days from the breakout of the shock wave at the stellar surface (see Figure 18)). | shell, increases by a factor $\sim$ 40 in $\sim$ 35 days from the breakout of the shock wave at the stellar surface (see Figure \ref{fig:modB_radius}) ). |
The expansion is nearly homologous, as witnessed by the nearly constant value of the photospheric velocity. | The expansion is nearly homologous, as witnessed by the nearly constant value of the photospheric velocity. |
At ~ 35 days, in the outermost layers the temperature reaches ~ 6000 K and hydrogen starts to recombine. | At $\sim$ 35 days, in the outermost layers the temperature reaches $\sim$ 6000 K and hydrogen starts to recombine. |
This marks the beginning of the recombination phase. | This marks the beginning of the recombination phase. |
During this phase, the evolution is characterized by the fast recombination of the outer zone, leading to the formation of a recombination wave (RW) that moves inward (in mass) like a flame. | During this phase, the evolution is characterized by the fast recombination of the outer zone, leading to the formation of a recombination wave (RW) that moves inward (in mass) like a flame. |
In the external layers, the temperature drops below ~ 6000 K and, consequently, the internal energy is efficiently radiated away because of the sudden decrease of the opacity (see Figures 13,, 15 and 16)). | In the external layers, the temperature drops below $\sim$ 6000 K and, consequently, the internal energy is efficiently radiated away because of the sudden decrease of the opacity (see Figures \ref{fig:modB_75d}, \ref{fig:modB_TvsR} and \ref{fig:modB_TvsM}) ). |
The RW marks the boundary between the inner envelope that is optically thick, ionized and hot (> 6000 K), and the outer layers that are optically thin, recombined and cooler (< 6000 K). | The RW marks the boundary between the inner envelope that is optically thick, ionized and hot $>$ 6000 K), and the outer layers that are optically thin, recombined and cooler $\lesssim$ 6000 K). |
The photosphere | The photosphere |
across (he eieht. CCDs of the 47 Tue field allowing a determination of the position of all stus in both our photometric and lishteurve database. | across the eight CCDs of the 47 Tuc field allowing a determination of the position of all stars in both our photometric and lightcurve database. |
We used a program to search ihe USNO CCD Astrograph Catalogue (UCACI) for astrometric standard stars within the field (D.P.Schmidt. 2003 private communication). | We used a program to search the USNO CCD Astrograph Catalogue (UCAC1) for astrometric standard stars within the field (B.P.Schmidt, 2003 private communication). |
Several hundred such stars were cross-identified. allowing for an accurate determination of the astrometric solution of our star lists. | Several hundred such stars were cross-identified, allowing for an accurate determination of the astrometric solution of our star lists. |
The rms residual of the astrometry was 720.15 aresecs. | The rms residual of the astrometry was $\sim$ 0.15 arcsecs. |
Our astronietry is presented in Fie... with the CCD number overplotted. | Our astrometry is presented in \ref{astromplot}, with the CCD number overplotted. |
The cluster core and our extensive coverage of the cluster is apparent. | The cluster core and our extensive coverage of the cluster is apparent. |
The V.V-I Colour Magnitude Diagram (CMD) used in the production of the variable colour data is presented in Fig.5.. | The V,V-I Colour Magnitude Diagram (CMD) used in the production of the variable colour data is presented in \ref{47TucCMD_total}. |
These data were originally taken at the MSSSO 40" telescope by Ixen Freeman ancl Michelle Doherty (o allow for a placement of any candidate transiting svstem on the V.V-I svstem. | These data were originally taken at the MSSSO $''$ telescope by Ken Freeman and Michelle Doherty to allow for a placement of any candidate transiting system on the V,V-I system. |
The data cover the same 5252’ FOV as the lighteurve dataset presented in Chis paper. ancl totalling 43.067 stars. | The data cover the same $\times$ $'$ FOV as the lightcurve dataset presented in this paper, and totalling 43,067 stars. |
The magnitude range is 12.021.0 in V and I. The CMD was calibrated against that presented by (1998). | The magnitude range is $\leqslant$ $\leqslant$ 21.0 in V and I. The CMD was calibrated against that presented by \citet{Kal98}. |
. The authors warn in that paper of svstematic errors caused by non-linearity of the OGLE CCD chip. | The authors warn in that paper of systematic errors caused by non-linearity of the OGLE CCD chip. |
For faint stars (hese errors are likely to be more significant. | For faint stars these errors are likely to be more significant. |
The OGLE dataset as available on their website (calibrated) was overplotted on top of our uncalibrated CMD data. | The OGLE dataset as available on their website (calibrated) was overplotted on top of our uncalibrated CMD data. |
Our data was then shifted until the two datasets overlay eachother. | Our data was then shifted until the two datasets overlay eachother. |
A histogram of the distribution of the stellar magnitudes was produced. and our data was shifted further until a more accurate match was found by comparing (he V magnitude of the Horizontal Branch. | A histogram of the distribution of the stellar magnitudes was produced, and our data was shifted further until a more accurate match was found by comparing the V magnitude of the Horizontal Branch. |
This was repeated for V-I calibration. | This was repeated for V-I calibration. |
This simple calibration method produced photometry accurate to <0.03 mag. across the full magnitude range. | This simple calibration method produced photometry accurate to $\leqslant$ 0.03 mag, across the full magnitude range. |
Fig.5 shows our CMD. | \ref{47TucCMD_total} shows our CMD. |
Variable stars in the V+R dataset were cross-identified by | Variable stars in the V+R dataset were cross-identified by |
studied we tried to combine Ca/LFe]. the ceccntricity and esiné in 3D-graph (see Figure. NJ). | studied we tried to combine [Ca/Fe], the eccentricity and $\vsini$ in 3D-graph (see Figure \ref{fevsini}) ). |
We obtained a weak negative correlation between the eccentricity ancl the projected. rotational velocity - faster rotating stars tended to have smaller eccentricities. | We obtained a weak negative correlation between the eccentricity and the projected rotational velocity - faster rotating stars tended to have smaller eccentricities. |
The dependence of the Am peculiarity on the elfective temperature is plotted on the Fig. 10.. | The dependence of the Am peculiarity on the effective temperature is plotted on the Fig. \ref{fteff}. |
There might be a trend of decreasing the peculiarity with the temperature - the correlation coefficient is. O.86-E0.06. | There might be a trend of decreasing the peculiarity with the temperature - the correlation coefficient is $\pm0.06$. |
We sought a possible dependence of the Xm peculiarity on the age and the mass of the stars. | We sought a possible dependence of the Am peculiarity on the age and the mass of the stars. |
With the exception of one star. 1196544. the other stars studied have very close ages - the cillerence between the voungest and the oldest star is 0.51 in logZ (where Tis the age in vears). | With the exception of one star, 196544, the other stars studied have very close ages - the difference between the youngest and the oldest star is 0.51 in $\log T$ (where $T$ is the age in years). |
The majority of the masses of the stars is also in very short interval - 0.5 in solar masses. | The majority of the masses of the stars is also in very short interval - 0.5 in solar masses. |
As a result from all our data we can not see any signs of dependence between these parameters. (mass and age) and the Am peculiarity. | As a result from all our data we can not see any signs of dependence between these parameters (mass and age) and the Am peculiarity. |
We also. studied the connection between the mücroturbulence and the. cllective temperature (νου lig. 113). | We also studied the connection between the microturbulence and the effective temperature (see Fig. \ref{fvturb}) ). |
There seems not to be any dependence between hese parameters for the temperature region 7000-9000Ix. But at higher temperatures it is possible to claim thatIx the microturbulence decreases with the increasing teniperature. | There seems not to be any dependence between these parameters for the temperature region K. But at higher temperatures it is possible to claim that the microturbulence decreases with the increasing temperature. |
These results fit relatively well with the conclusion given w Burkhart&Coupry(1992). | These results fit relatively well with the conclusion given by \citet{bc92}. |
.. Phe microturbulence is a pure fitting parameter which brings into agreement he abundances from weak and strong lines. | The microturbulence is a pure fitting parameter which brings into agreement the abundances from weak and strong lines. |
Lt does. not necessarily have the meaning of existing turbulent motions. | It does not necessarily have the meaning of existing turbulent motions. |
evertheless. this behaviour is in agreement with the expectations and with the fact that. supelicial convective zones ect thinner at higher temperatures ancl cease at about Kh. We also studied a possible dependence of the Am peculiarities on the microturbulenee but there is no clear correlation. | Nevertheless, this behaviour is in agreement with the expectations and with the fact that supeficial convective zones get thinner at higher temperatures and cease at about K. We also studied a possible dependence of the Am peculiarities on the microturbulence but there is no clear correlation. |
Apparently. the microturbulenee does not seem to destrov the Xam peculiarity. | Apparently, the microturbulence does not seem to destroy the Am peculiarity. |
We compared our measurements of the projected rotational velocities of the program stars to those obtained bv Abt&Morrell(1995). and by Roveretal.(2002). | We compared our measurements of the projected rotational velocities of the program stars to those obtained by \citet{am95} and by \citet{rgbgz02}. |
.. As it is seen from Figure 12. our measurements are in. good agreement with them. | As it is seen from Figure \ref{fvel} our measurements are in good agreement with them. |
The velocities of Abt&Morrell seers to be closer to our values while those of (2002) are often slightly higher. | The velocities of \citet{am95} seems to be closer to our values while those of \citet{rgbgz02} are often slightly higher. |
In this third. and last paper. we have analysed another six binaries from. our sample and derived their chemical composition. temperatures. gravities. projected. rotational velocities. masses and ages. | In this third and last paper, we have analysed another six binaries from our sample and derived their chemical composition, temperatures, gravities, projected rotational velocities, masses and ages. |
Phe obtained abundances of | The obtained abundances of |
the problem of observing less modes than predicted by the niocdels. | the problem of observing less modes than predicted by the models. |
Lt also provides a useful tool to determine the mass and width of the surface laver (see e.g. Brassareletal. 19091: ]xawaler&Bradle | It also provides a useful tool to determine the mass and width of the surface layer (see e.g. \citealt{brassard92}; ; \citealt{kawaler94}) ). |
y 1994). Charpinetetal.(2000). reported eravitv mode trapping due to the transition between the ielium core and the hydrogen-rich envelope in sdB models. | \citet{charpinet00} reported gravity mode trapping due to the transition between the helium core and the hydrogen-rich envelope in sdB models. |
They also identified: subtle. but. non-negligible departures rom uniform frequency spacing for pressure modes. caused ov the same chemical transition region. | They also identified subtle, but non-negligible departures from uniform frequency spacing for pressure modes, caused by the same chemical transition region. |
In sdD models. the maximum gradient in Brunt-Vaisalla (BV) frequency is produced at the transition of he helium radiative core to the hydrogen envelope. | In sdB models, the maximum gradient in Brunt-V\"aiis\"all\"a (BV) frequency is produced at the transition of the helium radiative core to the hydrogen envelope. |
In sd models the BY profile is more complex. since as evolution ooceeds. a C-O core builds up. resulting in two chemical ransitions: one from the C-O core to the He burning shell. and another one from the He radiative shell to a H1 burning shell. | In sdO models the BV profile is more complex, since as evolution proceeds, a C-O core builds up, resulting in two chemical transitions: one from the C-O core to the He burning shell, and another one from the He radiative shell to a H burning shell. |
Phis renders the mode trapping elfects more complex in scdO than in sdD mocels. | This renders the mode trapping effects more complex in sdO than in sdB models. |
We report in this study the. discovery of. g-moce trapping in a sdO model caused. by the ο. transition. | We report in this study the discovery of -mode trapping in a sdO model caused by the He/H transition. |
This provides a selection mechanism in the way that trapped mocles would be more easilv excited. | This provides a selection mechanism in the way that trapped modes would be more easily excited. |
Low- and intermediate racial order are found to be trapped by the deeper transition from the C-O/Le. but with no significant effects on the driving. | Low- and intermediate radial order are found to be trapped by the deeper transition from the C-O/He, but with no significant effects on the driving. |
The importance of gmocde trapping is evident as a way to probe the deep stellar interior. where gravity modes propagate. unattainable in any other wav: but also to derive information about the location and width of the chemical transitions and the mass of the envelopo. | The importance of -mode trapping is evident as a way to probe the deep stellar interior, where gravity modes propagate, unattainable in any other way; but also to derive information about the location and width of the chemical transitions and the mass of the envelope. |
The paper is organized as follows: in Section 2 we describe some general characteristics of the sdO mocel used. | The paper is organized as follows: in Section 2 we describe some general characteristics of the sdO model used. |
In Section 3 we describe the g-modoe trapping and the effects of cancelling out the LezL and €-Oο chemical transitions in BY frequenev on the trapping and driving of the modes. | In Section 3 we describe the -mode trapping and the effects of cancelling out the He/H and C-O/He chemical transitions in BV frequency on the trapping and driving of the modes. |
Section 4 does a similar treatment for. p-modoes. | Section 4 does a similar treatment for -modes. |
. Finally. Section 5 presents a summary and Section 6 the discussion and conclusions. | Finally, Section 5 presents a summary and Section 6 the discussion and conclusions. |
This theoretical exercise. arose as an exploration of a particular sd model. whose properties. were thoroughly described. in Paper ancl parameters given in ‘Table 1.. | This theoretical exercise arose as an exploration of a particular sdO model, whose properties were thoroughly described in Paper and parameters given in Table \ref{tab:model}. |
The model. built with code (Lawlor&MacDonald 2006).. comes from a 1 M7 star on the pre-main sequence with enhanced mass loss rate on the red giant branch. that drives the star to evolve to higher temperatures at constant uminosity. | The model, built with code \citep{lawlor06}, comes from a 1 $\odot$ star on the pre-main sequence with enhanced mass loss rate on the red giant branch, that drives the star to evolve to higher temperatures at constant luminosity. |
A delaved helium-IDash puts the star back in the iorizontal. branch. and further evolution brings the star to he sdO domain. | A delayed helium-flash puts the star back in the horizontal branch, and further evolution brings the star to the sdO domain. |
The sdO models have developed a carbon-oxvgen core while helium and hydrogen shell burning is still woduced (Fig. | The sdO models have developed a carbon-oxygen core while helium and hydrogen shell burning is still produced (Fig. |
E. Heft). | \ref{fig:phys2epsn2} left). |
Phe change in chemical composition rom the €-O core to the He burning shell. and. [rom the Le radiative laver to the HE burning shell is seen in the two steep peaks in BY frequency and subtler transitions in Lanib requeney (Fig. | The change in chemical composition from the C-O core to the He burning shell, and from the He radiative layer to the H burning shell is seen in the two steep peaks in BV frequency and subtler transitions in Lamb frequency (Fig. |
1. right) | \ref{fig:phys2epsn2} right). |
The model. which was subject to a non adiabatic analysis with (Alovaetal.2004: Mova&Carrico 2008)). was found stable in the frequency. range ~0.5 to 25 mllz. | The model, which was subject to a non adiabatic analysis with \citealt{moya04}; \citealt{moya08}) ), was found stable in the frequency range $\sim$ 0.5 to 25 mHz. |
However. an oscillatory behaviour of the growth rate at low ancl intermediate frequencies. (see fig. | However, an oscillatory behaviour of the growth rate at low and intermediate frequencies (see fig. |
10 of Paper Ll) caught our attention. | 10 of Paper I) caught our attention. |
That drove us to explore the behaviour of the kinetic energy of the modes. which revealed that the model experienced. mode trapping ellects. | That drove us to explore the behaviour of the kinetic energy of the modes, which revealed that the model experienced mode trapping effects. |
The logarithm of the total kinetic energy. of each g2mocde | The logarithm of the total kinetic energy of each -mode |
in PNSs cannot be strong sources of gravitational wave radiation. | in PNSs cannot be strong sources of gravitational wave radiation. |
Lastly. future core-collapse simulations must account for sub-grid mode coupling in order to suppress the artificial growth of low-order core oscillations in cases where such growth inlluences the simulation results. | Lastly, future core-collapse simulations must account for sub-grid mode coupling in order to suppress the artificial growth of low-order core oscillations in cases where such growth influences the simulation results. |
We thank A. Burrows for providing the PNS models and for helpful comments on an earlier draft of this paper. P. Arras for a valuable correspondence on the caleulation of coupling coefficients. and T. Fhompson. T. van Hoolst. and Y. Wu for helpful suggestions and input. | We thank A. Burrows for providing the PNS models and for helpful comments on an earlier draft of this paper, P. Arras for a valuable correspondence on the calculation of coupling coefficients, and T. Thompson, T. van Hoolst, and Y. Wu for helpful suggestions and input. |
This work was supported by the “Pheoretical Astrophysies Center at UC Berkeley and by NASA grant NNCGOGCGIG6SCG and the David Lucile Packard Founcation. | This work was supported by the Theoretical Astrophysics Center at UC Berkeley and by NASA grant NNG06GI68G and the David Lucile Packard Foundation. |
The classical double. ΕΠΗ radio sources are among the most luminous extragalactic racio sources (Fanarolf Riley 1974). | The classical double FRII radio sources are among the most luminous extragalactic radio sources (Fanaroff Riley 1974). |
They are characterised by two steep spectrum radio lobes. symmetrically disposed. with respect. το the host galaxy or quasar. and. according to the standard. picture. these are powered. by beams or jets originating in an active galactic nucleus. | They are characterised by two steep spectrum radio lobes, symmetrically disposed with respect to the host galaxy or quasar, and, according to the standard picture, these are powered by beams or jets originating in an active galactic nucleus. |
The identification of Ilat-spectrum central radio cores in most of these sources has enabled a number of analyses of the kinematics of their lobes and hot-spots to be undertaken (Longair Riley 1979. Zieba ChDyvzv 1991. Best 1995). | The identification of flat-spectrum central radio cores in most of these sources has enabled a number of analyses of the kinematics of their lobes and hot-spots to be undertaken (Longair Riley 1979, Zieba Chyzy 1991, Best 1995). |
Structural asvmimetries and misalignments of the hot-spots ancl radio lobes have been used to investigate their intrinsic. properties ancl to test unification schemes for racio quasars and radio galaxies (Scheuer LOST: Barthel 1987. 1989). | Structural asymmetries and misalignments of the hot-spots and radio lobes have been used to investigate their intrinsic properties and to test unification schemes for radio quasars and radio galaxies (Scheuer 1987; Barthel 1987, 1989). |
In the simplest picture. it was conjectured that the radio lobes ancl hot-spots moved: out svnunetrically [rom an active nucleus at a significantly relativistic speed and the observed. structural asvmmetries could then be attributed to the dillerences in light travel times from the lobes to the observer (livle Longair 1967). | In the simplest picture, it was conjectured that the radio lobes and hot-spots moved out symmetrically from an active nucleus at a significantly relativistic speed and the observed structural asymmetries could then be attributed to the differences in light travel times from the lobes to the observer (Ryle Longair 1967). |
Many. studies have been mace of the probability distribution of the velocities of the radio source components from the observed. distributions of the ratio of corehot-spot distances (Longair Riley 1979: ]xatgert-Meikelijn 1980: Danhatti 1980: Best 1995). | Many studies have been made of the probability distribution of the velocities of the radio source components from the observed distributions of the ratio of core–hot-spot distances (Longair Riley 1979; Katgert-Meikelijn 1980; Banhatti 1980; Best 1995). |
Phe mean velocity of advance of the hot-spots was found to be z0.2e. with a considerable spread. about the mean velocity. some values greater than 04e being found. | The mean velocity of advance of the hot-spots was found to be $\geq 0.2c$, with a considerable spread about the mean velocity, some values greater than $0.4c$ being found. |
There are however significant problems with this moclel. | There are however significant problems with this model. |
MeCarthy (1991) studied: the spatial distribution of thermal emission-line gas about. powerful. FRIL radio sources and found that the asvmmetry in the distribution of the ionisccd gas was correlated with the structural asvinmetry of the radio lobes. strongly suggesting. that environmental asvmimetries play a significant rolle. | McCarthy (1991) studied the spatial distribution of thermal emission-line gas about powerful FRII radio sources and found that the asymmetry in the distribution of the ionised gas was correlated with the structural asymmetry of the radio lobes, strongly suggesting that environmental asymmetries play a significant rôlle. |
“Phis correlation between optical ancl radio asvnimietrics can be naturally attributed to. clumpy environmental. effects (οον 1989a. b). | This correlation between optical and radio asymmetries can be naturally attributed to clumpy environmental effects (Pedelty 1989a, b). |
A further possibility is that the jets of powerful FRIL sources might well be intrinsically asvmumetric. | A further possibility is that the jets of powerful FRII sources might well be intrinsically asymmetric. |
The analysis of Wardle Aaron (1997). of | The analysis of Wardle Aaron (1997) of |
The cosmological relevance of local starbursts is due to their similarities with their analogs at lieh redshift. | The cosmological relevance of local starbursts is due to their similarities with their analogs at high redshift. |
They may therefore provide extremely valuable laboratories to study with hieh resolution and seusitivitv relevant physical processes associated with the starburst activity iu imclh better detail. | They may therefore provide extremely valuable laboratories to study with high resolution and sensitivity relevant physical processes associated with the starburst activity in much better detail. |
One of the most suitable euvironniecnts for these purposes can boe the intrieuiug class of very compact. extremely forming galaxy (SFC) at low redshift (0.11.<:oOX (035) recently discovered by volunteers inthe "Galaxy Zoo” project (Lintottetal. 2008). | One of the most suitable environments for these purposes can be the intriguing class of very compact, extremely star-forming galaxy (SFG) at low redshift $0.11 \la z
\la 0.35$ ) recently discovered by volunteers inthe “Galaxy Zoo” project \citep{Lintott}. |
. Those ealaxies. popularly referred to as vereen peas” (hereafter CPs). were first reported and studied at some length by Cordanmoneet CO9).. who classified more | These galaxies, popularly referred to as “green peas” (hereafter GPs), were first reported and studied at some length by \citet[][hereafter C09]{C09}, , who classified more |
GeV. the cascade radiation for the two models is roughly comparable. though there is more radiation at the lowest chereics for the Stecker model. | GeV, the cascade radiation for the two models is roughly comparable, though there is more radiation at the lowest energies for the Stecker model. |
Above 10 GeV iu the PLE model of blazars. there is wore cascade radiation for the Ciüliiore model than for the Stecker model due to the chhanced absorption in the Stecker model. | Above $10$ GeV in the PLE model of blazars, there is more cascade radiation for the Gilmore model than for the Stecker model due to the enhanced absorption in the Stecker model. |
As a result. the spectra at high energies are vastly different. | As a result, the spectra at high energies are vastly different. |
As noted in ?.. such au absorption feature is much more proninent iu he Secker case than in the Cülinore case owjug to the lugher UV backerouud iu the Stecker EBL model. | As noted in \citet{ven09}, such an absorption feature is much more prominent in the Stecker case than in the Gilmore case owing to the higher UV background in the Stecker EBL model. |
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