ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

We also improved the telescope guiding system's stability.

This led to an improvement in the precision of the light curves which is evident in the latest light curve of WASP- taken after the upgrades (see Fig.

This led to an improvement in the precision of the light curves which is evident in the latest light curve of WASP-21 taken after the upgrades (see Fig.

1).

In the November 2010 observations the variation in position is less than 2 pixels in the or direction and 4 pixels in the y.

In the November 2010 observations the variation in position is less than 2 pixels in the $x$ direction and 4 pixels in the $y$.

An accurate estimate of the photometric errors is important to obtain reliable system. parameters.

An accurate estimate of the photometric errors is important to obtain reliable system parameters.

Our first estimate of errors or each light curve includes only the shot noise. readout and background noise. which underestimates the true errors.

Our first estimate of errors for each light curve includes only the shot noise, readout and background noise, which underestimates the true errors.

To obtain a more reliable estimate we begin by scaling the errors of each ight curve so that the reduced X7 of the best fitting model is 1.0.

To obtain a more reliable estimate we begin by scaling the errors of each light curve so that the reduced $\chi^2$ of the best fitting model is 1.0.

This resulted in the multiplication of the errors by 1.97. 1.22 and 44. for the 2008. 2009 and 2010 light curves. respectively.

This resulted in the multiplication of the errors by $1.97$ , $1.22$ and $1.44$, for the 2008, 2009 and 2010 light curves, respectively.

We jen. calculated the time-correlated noise following the procedure rom ?..

We then calculated the time-correlated noise following the procedure from \citet{gillon2009}.

Using the residuals of the best fit model. we estimated ye amplitude of the red noise. σι. to be 150 ppm. 250 ppm and 50 ppm. for the 2008. 2009 and 2010 light curves. respectively.

Using the residuals of the best fit model, we estimated the amplitude of the red noise, $\sigma_r$ to be $150\,$ ppm, $250\,$ ppm and $150\,$ ppm, for the 2008, 2009 and 2010 light curves, respectively.

These were added in quadrature to the rescaled photometric errors and were used in the final Markov Chain Monte Carlo (MCMC) chains.

However. ? found that this "time-averaging" method of estimating the correlated noise can still underestimate the uncertainties by 15-30 per cent.

However, \citet{carter2009} found that this “time-averaging” method of estimating the correlated noise can still underestimate the uncertainties by 15-30 per cent.

To determine the planetary and orbital parameters. we fitted the three RISE light curves of WASP-21b simultaneously.

To determine the planetary and orbital parameters, we fitted the three RISE light curves of WASP-21b simultaneously.

We used the ?. transit model parametrised by the normalised separation of the planet. @//?,.. ratio of planet radius to star radius. £2),/22). orbital inclination. 7. and the transit epoch. 76. of each light curve.

We used the \citet{Mandel2002} transit model parametrised by the normalised separation of the planet, $a/R_*$, ratio of planet radius to star radius, $ R_p/R_* $, orbital inclination, $i$, and the transit epoch, $T_0$, of each light curve.

Our model was originally. developed © measure transit timing variations of exoplanets.

Our model was originally developed to measure transit timing variations of exoplanets.

Following ? qat found no evidence for a significant orbital eccentricity of WASP-21b we adopt a circular orbit.

Following \citet{Bouchy2010} that found no evidence for a significant orbital eccentricity of WASP-21b we adopt a circular orbit.

We included the quadratic limb diirkening (LD) coefficients for the RISE filter V+R from the mocels of α=0.45451 and b=0.21017

We included the quadratic limb darkening (LD) coefficients for the RISE filter V+R from the models of \citet{Howarth2010}: $a=0.45451$ and $b=0.210172$.

These were calculated for ει= 5800K. g=44.2 and [M/H] 0.5 to match 1e Stellar parameters from ?..

These were calculated for $T_{eff} = 5800\,$ K, 4.2 and [M/H] $-$ 0.5 to match the stellar parameters from \citet{Bouchy2010}.

We initially kept the limb darkening parameters fixed during the fit.

For each light curve. we included two extra parameters to account for a linear normalization.

For each light curve, we included two extra parameters to account for a linear normalization.

Therefore. 12. parameters were fitted.

Therefore, 12 parameters were fitted.

Besides the linear normalization. no extra trends were removed fromthe light curve.

Besides the linear normalization, no extra trends were removed fromthe light curve.

To obtain the best tit parameters and uncertainties. we used

To obtain the best fit parameters and uncertainties, we used

Based on equation (1). the investigation of the EAS inverse problem. which is the reconstruction of cherey spectra of primary nuclei by the observable EAS electron and muon size spectra at observation level males sense ouly using given functions for Πο primary energy spectra with given uuknown spectral parameters (so called parameterization of equation (1)) as it is done in[S.18.19. 201.

Based on equation (1), the investigation of the EAS inverse problem, which is the reconstruction of energy spectra of primary nuclei by the observable EAS electron and muon size spectra at observation level makes sense only using given functions for unknown primary energy spectra with given unknown spectral parameters (so called parameterization of equation (1)) as it is done in\cite{Glass,STPB,TH,TB}. .

(2010).

.

. In a recent paper, analysed the multiwavelength data from the 2007-2008 WEBT campaign, including three pointings by XMM-Newton.

In a recent paper, analysed the multiwavelength data from the 2007–2008 WEBT campaign, including three pointings by XMM-Newton.

The XMM-Newton data revealed a UV excess, which was interpreted to be due to thermal emission from the accretion disc, as well as a spectral curvature in the X-ray band.

The authors constructed spectral energy distributions (SEDs) of BL Lacertae corresponding to various epochs where the source was in different brightness states, using both their own data and data from the literature.

They applied the inhomogeneous, rotating helical jet model by 1999,,2003,, 2004)mainBodyCitationEnd220] to fit the SEDs, and suggested that the broad-band spectral properties of BL Lacertae may result from the combination of two synchrotron emission components with their self inverse-Compton emission, plus a thermal component from the disc.

They applied the inhomogeneous, rotating helical jet model by , ] to fit the SEDs, and suggested that the broad-band spectral properties of BL Lacertae may result from the combination of two synchrotron emission components with their self inverse-Compton emission, plus a thermal component from the disc.

Subsequently, analysed optical spectra acquired in the same period with the 3.56 m Telescopio Nazionale Galileo (TNG).

They found a broad Ha emission line, with luminosity of ~4x10*'ergs! and FWHM of ~ 4600kms'!, even brighter than that found in 1997 by and2000).

They found a broad $\alpha$ emission line, with luminosity of $\sim 4 \times 10^{41} \, \rm erg \, s^{-1}$ and FWHM of $\sim 4600 \rm \, km \, s^{-1}$ , even brighter than that found in 1995--1997 by and.

. This favours the hypothesis that the UV excess is caused by thermal emission from the accretion disc, the most likely source of ionising photons for the broad line region.

This favours the hypothesis that the UV excess is caused by thermal emission from the accretion disc, the most likely source of ionising photons for the broad line region.

The multiwavelength data available for the analysis lacked simultaneous information in the y- band, so that the inverse-Compton spectral region was poorly constrained.

The multiwavelength data available for the analysis lacked simultaneous information in the $\gamma$ -ray band, so that the inverse-Compton spectral region was poorly constrained.

But in 2008 the Fermi satellite was able to detect BL Lacertae2010a),, even if in a low state compared to the past detections by the Compton Gamma Ray (CGRO,1999,, 1997)).

But in 2008 the Fermi satellite was able to detect BL Lacertae, even if in a low state compared to the past detections by the Compton Gamma Ray (CGRO, ).

In the same period, observations in the UV and bands were performed by Swift, while in the optical, near- mm and cm radio bands the source was monitored by the GLAST-AGILE Support Program (GASP) of the WEBT.

In the same period, observations in the UV and X-ray bands were performed by Swift, while in the optical, near-IR, mm and cm radio bands the source was monitored by the GLAST-AGILE Support Program (GASP) of the WEBT.

This offered the unique opportunity to study the source emission over a very extended spectral range.

The results of this further investigation effort on BL Lacertae are presented in this paper.

The GASP was born in 2007 as a WEBT project, with the aim of monitoring a list of 28 y-ray loud blazars in the optical, near- mm, and cm radio bands during the y-ray observations of the and (formerly GLAST) satellites2009a).

The GASP was born in 2007 as a WEBT project, with the aim of monitoring a list of 28 $\gamma$ -ray loud blazars in the optical, near-IR, mm, and cm radio bands during the $\gamma$ -ray observations of the and (formerly GLAST) satellites.

. Data are collected periodically by the WEBT President, who checks the consistency of the various datasets.

Data are collected periodically by the WEBT President, who checks the consistency of the various datasets.

The GASP light curves are then available for multiwavelength studies, mostly in the framework of the GASP collaboration with the AGILE and Fermi research teams.

The GASP data presented in this paper were taken at the observatories listed in Table 1..

The GASP data presented in this paper were taken at the observatories listed in Table \ref{obs}.

The optical data were calibrated with respect to a common choice of reference stars in the same field of the source in U and B bands; in V, R, and J).

The optical data were calibrated with respect to a common choice of reference stars in the same field of the source in $U$ and $B$ bands; in $V$, $R$, and $I$ ).

The source photometry was evaluated from a circular region with an 8 arcsec aperture radius, while the background was taken in a surrounding annulus with 10 and 16 arcsec radii.

In this way the measureis essentially seeing-independent and all datasets are affected by

are operated on the island of La Palma by. the Isaac Newton Group in the Spanish Observatorio cel Roque de los) Muchachos of the Instituto. de Astrolisica cde Canarias.

are operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias.

This research. has mace use of the NASA/IPAC Extragalactic Database (NIZD) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Acronautics ancl Space Administration.

This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.

The authors thank Neville Shane for both helping out with the observations and calculations provided in this paper.

We also thank the referee for comments which improved the content of this paper.

our study.

This comparison clearly shows the strong correlation between high mass loss rates and variability.

Mauas et al. (2006))

Mauas et al. \cite{Mauas06}) )

derive a mass loss rate of a few 107? Μο on the upper part of the giant branch of NGC 2808 from the analysis of chromospheric lines.

derive a mass loss rate of a few $^{-9}$ $M_{\odot}$ on the upper part of the giant branch of NGC 2808 from the analysis of chromospheric lines.

Ho line components at an outflow velocity were found in a large sample of red giants by Cacciari et al. (2004)).

$\alpha$ line components at an outflow velocity were found in a large sample of red giants by Cacciari et al. \cite{Cacciari04}) ).

A Spitzer study of this cluster has been obtained already (Fabbri et al. 2008)),

A Spitzer study of this cluster has been obtained already (Fabbri et al. \cite{Fabbri08}) ),

but the results are not available yet.

From the horizontal branch morphology of this cluster D'Antona Caloi (2008)) and Dalessandro et al. (2011))

From the horizontal branch morphology of this cluster D'Antona Caloi \cite{dAC08}) ) and Dalessandro et al. \cite{dal10}) )

derive an average mass lost before the red clump phase of Me.

derive an average mass lost before the red clump phase of $M_{\odot}$.

We have shown that modelling of the LPVs in 47 Tuc allows one to derive theamount of mass lost on the RGB (Lebzelter Wood 2005)).

We have shown that modelling of the LPVs in 47 Tuc allows one to derive theamount of mass lost on the RGB (Lebzelter Wood \cite{LW05}) ).

It is our aim to carry out the same pulsation analysis for NGC 362 and NGC 2808.

In addition, given the large range in the helium abundance suspected in NGC 2808, we aim to see how the helium abundance affects the LPV pulsation periods.

In particular, we aim to see if the position on the period-luminosity diagram can give a clue to the helium abundance.

Based on the summary of abundance measurements in the two clusters detailed in the Introduction, we assume a metal abundance Z=0.001 for both clusters.

Models have been made with helium mass fractions Y=0.25, 0.3 and 0.4.

Note that these are, respectively, the estimated helium mass fractions given by Dalessandro et al. (2011))

Note that these are, respectively, the estimated helium mass fractions given by Dalessandro et al. \cite{dal10}) )

for red horizontal branch, blue horizontal branch and extreme horizontal branch stars in NGC 2808.

Mass loss was added to some models assuming a Reimers’ Law (Reimers 1975)).

Mass loss was added to some models assuming a Reimers' Law (Reimers \cite{r75}) ).

The Reimers mass loss rate was multiplied by a factor 7=0.4.

The Reimers mass loss rate was multiplied by a factor $\eta$ =0.4.

This factor was chosen as it leads to a termination of the AGB near the maximum observed luminosity of the variable stars in NGC 362 and NGC 2808.

It also closely reproduces the masses estimated for the horizontal branch stars in the two clusters (see below), and it is close to the value of 0.33 we found to apply in 47 Tuc.

In order to compute the declining stellar mass along the RGB and AGB in the presence of mass loss, evolution rates as a function of luminosity were obtained from the evolution tracks of Bertelli et al (2008)).

In order to compute the declining stellar mass along the RGB and AGB in the presence of mass loss, evolution rates as a function of luminosity were obtained from the evolution tracks of Bertelli et al \cite{Bertelli08}) ).

Teg in our calculations was forced to be the same as the mean observed Τεῃ on the giant branch (see Fig.

$T_{\rm eff}$ in our calculations was forced to be the same as the mean observed $T_{\rm eff}$ on the giant branch (see Fig.

9 or Fig. 11))

\ref{hrd_mconst} or Fig. \ref{hrd_eta0.4}) )

so that the radius used in the Reimers’ Law formula was appropriate for NGC 362 and NGC 2808.

so that the radius used in the Reimers' Law formula was appropriate for NGC 362 and NGC 2808.

Following the discussion in the Introduction, we have assumed an age of 10x10? years for both clusters and for each population of different helium abundance.

Following the discussion in the Introduction, we have assumed an age of $\times 10^{9}$ years for both clusters and for each population of different helium abundance.

The tracks and isochrones of Bertelli et al. (2008))

The tracks and isochrones of Bertelli et al. \cite{Bertelli08}) )

were used to obtain the initial mass appropriate for giant branch stars of this age, Y and Z. Initial masses of 0.86, 0.79 and 0.65 Μο were found for helium abundances of 0.25, 0.3 and 0.4, respectively.

were used to obtain the initial mass appropriate for giant branch stars of this age, Y and Z. Initial masses of 0.86, 0.79 and 0.65 $M_{\odot}$ were found for helium abundances of 0.25, 0.3 and 0.4, respectively.

Linear pulsation models were computed with the code described in Lebzelter Wood (2005)).

Linear pulsation models were computed with the code described in Lebzelter Wood \cite{LW05}) ).

A mixing length of 1.7 pressure scale heights was used in all models as this was found to reproduce the observed Τεῃ of the giant branch well.

A mixing length of 1.7 pressure scale heights was used in all models as this was found to reproduce the observed $T_{\rm eff}$ of the giant branch well.

Models without mass loss are shown in the Hertzsprung- in Fig. 9..

Models without mass loss are shown in the Hertzsprung-Russell-diagram in Fig. \ref{hrd_mconst}. .

The My. and Teg values of the

The $M_{\rm bol}$ and $T_{\rm eff}$ values of the

at magnetic poles rather than at stellar equator.

This completely changes the topology of accretion How but the major conchisious about the effect ou stellar cooling are likely to hold.

This completely changes the topology of accretion flow but the major conclusions about the effect on stellar cooling are likely to hold.

Indeed. the magnetically channeled eas travels towards the stellar surface at a &ood fraction of the free-fall velocity ancl at some point it nist pass through the radiative shock. after which it acciunulates at the top of the maeuetospheric colunn of accreted material. as schematically indicated in Figure 6..

Indeed, the magnetically channeled gas travels towards the stellar surface at a good fraction of the free-fall velocity and at some point it must pass through the radiative shock, after which it accumulates at the top of the magnetospheric column of accreted material, as schematically indicated in Figure \ref{fig:mag_pole}.

Total energv release within the shock aud maguetospheric column is comparable to that occurring if the accretion disk were extending all the wav to the stellar surface.

Total energy release within the shock and magnetospheric column is comparable to that occurring if the accretion disk were extending all the way to the stellar surface.

This hot cohuun of accreted material illuniuates the surface of the star leading to the same suppression of intrinsic stellar τις as we discussed in this work.

This hot column of accreted material illuminates the surface of the star leading to the same suppression of intrinsic stellar flux as we discussed in this work.

In this case. however. irradiation 1s strongest near the magnetic poles while the magnetic equator is likely to be the coolest part of the stellarsurface’.

In this case, however, irradiation is strongest near the magnetic poles while the magnetic equator is likely to be the coolest part of the stellar.

.. Caleulatiou of stellar radiation and mteerated huuinositv iu this case would involve coustructing a model for the maguetosphieric cohuun structure and its radiative propertics.

Calculation of stellar irradiation and integrated luminosity in this case would involve constructing a model for the magnetospheric column structure and its radiative properties.

The impact of these details on the structure and evolution of young stars should be addressed by. future work.

The impact of these details on the structure and evolution of young stars should be addressed by future work.

Liuninosity of voung stars actively accreting from the civetuustellar disk can be siguificautlv affected bv the radiation which is produced in the inner parts of the disk ancl is intercepted by the stellar surface.

Luminosity of young stars actively accreting from the circumstellar disk can be significantly affected by the radiation which is produced in the inner parts of the disk and is intercepted by the stellar surface.

We showed that if a star gaius its mass via disk accretion on timescale of several 10° vr then the radiative flux caused by viscous dissipation in the disk is more than sufiicicut to increase the surface temperature of the star above the photospheric temperature that an isolated star with the same mass and radius would lave.

We showed that if a star gains its mass via disk accretion on timescale of several $10^5$ yr then the radiative flux caused by viscous dissipation in the disk is more than sufficient to increase the surface temperature of the star above the photospheric temperature that an isolated star with the same mass and radius would have.

πασάο by the disk is strongest iu the equatorial regions aud is almost neelieible near the poles.

Irradiation by the disk is strongest in the equatorial regions and is almost negligible near the poles.

Au outer radiative zone of almost constant temperature forms above the fully convective iuterior iu the strongly iradiated parts of he stellar surface.

An outer radiative zone of almost constant temperature forms above the fully convective interior in the strongly irradiated parts of the stellar surface.

This leads to the local suppression of iutrinsic enerev flux escaping from the stellar iuterior.

This leads to the local suppression of intrinsic energy flux escaping from the stellar interior.

We have demonstrated that there are two distinct nodes in which a fully convective object. can cool: nainly through the cool high-latitude polar regious or oxedonuünautlv through the low-latitude parts of the stellar surface.

We have demonstrated that there are two distinct modes in which a fully convective object can cool: mainly through the cool high-latitude polar regions or predominantly through the low-latitude parts of the stellar surface.

A particular regime of cooling iu a given object is set by the opacity behavior aud the adiabatic cluperature eradicut μμ iu the outer radiative zone.

A particular regime of cooling in a given object is set by the opacity behavior and the adiabatic temperature gradient $\nabla_{ad}$ in the outer radiative zone.

Accreting voung stars and brown dwirfs cool mainly hrough the polar regious while forming giaut planets cool through the whole surface.

Accreting young stars and brown dwarfs cool mainly through the polar regions while forming giant planets cool through the whole surface.

luteerated stellar Lhuninositv iu accreting case Is suppressed compared to the case of an isolated object. by up to a factor of several in some classes. of objects (actively accreting brown dwarts aud planets stars forming in gravitationally unstable disks in galactic unclei).

Integrated stellar luminosity in accreting case is suppressed compared to the case of an isolated object, by up to a factor of several in some classes of objects (actively accreting brown dwarfs and planets, stars forming in gravitationally unstable disks in galactic nuclei).

This leads to larger radii of bracdiated objects aud ay affect the initial conditions which are used to calculate the evolution of the low-mass objects ou timescales of ~10 Myr after their formation.

This leads to larger radii of irradiated objects and may affect the initial conditions which are used to calculate the evolution of the low-mass objects on timescales of $\sim 10$ Myr after their formation.

Existence of external radiative zoue may facilitate retention of dust in the atmospheres of brown chwarfs aud planets. aud may affect the streneth of magnetic Bold generated by internal dynamo in convective objects.

Existence of external radiative zone may facilitate retention of dust in the atmospheres of brown dwarfs and planets, and may affect the strength of magnetic field generated by internal dynamo in convective objects.

Some of the results obtained iu this work iav be yplicable to accreting white dwarf and neutron star SVSTCLUS.

Some of the results obtained in this work may be applicable to accreting white dwarf and neutron star systems.