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Flat-field frames were obtained for cach filter by subtracting he median dark frame frou the sky rales. normalizing the resulting frames and calculating 1ο media. | Flat-field frames were obtained for each filter by subtracting the median dark frame from the sky frames, normalizing the resulting frames and calculating the median. |
We then subtracted the dark frame frou i6. object παλιο, divided by the normalized flat-field ud subtracted the correspoucding sky frame. | We then subtracted the dark frame from the object frame, divided by the normalized flat-field and subtracted the corresponding sky frame. |
Sometimes. nother constant had to be subtracted in order to attain 1e zero backeround level. | Sometimes, another constant had to be subtracted in order to attain the zero background level. |
In order to align the resulting rales to better than a pixel fraction we used the non-saturated stars in these frames or. when there were uo stars. we used the centre of the ealaxy. | In order to align the resulting frames to better than a pixel fraction we used the non-saturated stars in these frames or, when there were no stars, we used the centre of the galaxy. |
For the fiux calibration a nuniber of standard stars were observed and mean extinctioncoefficients were applied (sy and Say dn magnitudes 1 are respectively 0.12. aud 08 for Las Campanas aud La Silla. 0.12 and 0.10. for he TDBL. 0.057 aud 0.012 for the NOT aud ky-=0.1 ‘or Calar Alto) | For the flux calibration a number of standard stars were observed and mean extinctioncoefficients were applied $_J$ and $_{K'}$ in magnitudes $^{-1}$ are respectively 0.12 and 0.08 for Las Campanas and La Silla, 0.12 and 0.10 for the TBL, 0.057 and 0.042 for the NOT and $_{K'}$ =0.1 for Calar Alto). |
The error level ou the standard stars amount toLO%. | The error level on the standard stars amount to. |
. In cohuun 7 of Table 2 we eive the isophotal farcsecy corresponding to 20 of the ckeround. | In column 7 of Table \ref{journal} we give the isophotal $^2$ corresponding to $\sigma$ of the background. |
We have retrieved UST calibrated infrared miasges obtained with NICMOS (filter FLGOW) for the cielit ealaxies for which such data was available (see column ] in Table 2)). | We have retrieved HST calibrated infrared images obtained with NICMOS (filter F160W) for the eight galaxies for which such data was available (see column 1 in Table \ref{journal}) ). |
Iv tages of the galaxies iu erev scale with isocontours overlaid ave shown if in Figs. | K' images of the galaxies in grey scale with isocontours overlaid are shown if in Figs. |
la-2sa. | 1a-28a. |
All figures are only available in electronic Orn. | All figures are only available in electronic form. |
Iu order to detect he preseuce of features showing a departure from radial sauiuetry such as bars aud spira arnas. we have used a iaskine technique. | In order to detect the presence of features showing a departure from radial symmetry such as bars and spiral arms, we have used a masking technique. |
We fltere the original images with a box of 2-3 times the FWIIN of the secine Guecdian filter) and divided the observe nuage bv the filtered onc. | We filtered the original images with a box of 2-3 times the FWHM of the seeing (median filter) and divided the observed image by the filtered one. |
The resulting images (hereafter called the sharp-divided nuages to clifferenciate them from those obtained by subtraction. ic. the so-callec uusharp masking techuique) are shown iu Figs. | The resulting images (hereafter called the sharp-divided images to differenciate them from those obtained by subtraction, i.e. the so-called unsharp masking technique) are shown in Figs. |
11-25). | 1b-28b. |
This techuique is very well suited to trace asviunetries 1u the light distribution. such as bars. spiral aris. dust lanes. vines: df allows the subtraction of the diffuse background In a very convenient wav to look for subtle. small-scale variations and diseuss the possible presence of both dust extiuguished aud more buninous regious (Sofie et al. | This technique is very well suited to trace asymmetries in the light distribution, such as bars, spiral arms, dust lanes, rings; it allows the subtraction of the diffuse background in a very convenient way to look for subtle, small-scale variations and discuss the possible presence of both dust extinguished and more luminous regions (Sofue et al. |
199|: Maarquez Moles 1996: Márquez et al. | 1994; Márrquez Moles 1996; Márrquez et al. |
1996: Erwin Sparke 1999: Laine ct al. | 1996; Erwin Sparke 1999; Laine et al. |
1999). | 1999). |
The preseuce of bars can be determined quantitatively by studving the behaviour of isophotal position angles (PAs) and ellipticifies (es): a bar is characterized by a local maxuuunmn in the e correspouding to a constant PA (Wozniak et al. | The presence of bars can be determined quantitatively by studying the behaviour of isophotal position angles (PAs) and ellipticities $\epsilon$ s): a bar is characterized by a local maximum in the $\epsilon$ corresponding to a constant PA (Wozniak et al. |
1995: Friedli ct al. | 1995; Friedli et al. |
1996: Jungwiert et al | 1996; Jungwiert et al. |
1997). | 1997). |
Therefore. we have ft ellipses to the isophotes (Jedrzejewski1987) with the IRΔΕ ellipse task in stsdas.analvsis.isophote. allowing the ceuter. PA aud € to vary frou oue isophote to the next. | Therefore, we have fit ellipses to the isophotes (Jedrzejewski1987) with the IRAF ellipse task in stsdas.analysis.isophote, allowing the center, PA and $\epsilon$ to vary from one isophote to the next. |
Foreeround stars have been masked out. | Foreground stars have been masked out. |
The resulting paraiecters allow the reconstruction of a model galaxy. | The resulting parameters allow the reconstruction of a model galaxy. |
The difference )oetween the original image aud the model (hereafter callec the ffference mages) are particularly useful to trace all the features that cannot be described with ellipses. as spiral arnis and boxy or peauut-Iike components. | The difference between the original image and the model (hereafter called the fference images) are particularly useful to trace all the features that cannot be described with ellipses, as spiral arms and boxy or peanut-like components. |
They are given in Fiex. | They are given in Figs. |
Le-28c and discussed separately for cach galaxy. | 1c-28c and discussed separately for each galaxy. |
When the PA suddenly changes from one isophote to the next. the model cannot work for the regious left inside (see Figs. | When the PA suddenly changes from one isophote to the next, the model cannot work for the regions left inside (see Figs. |
lc. 8c. 9c. 11ο. 12ο. 15ο, 18e. 19ο, 25e and 26c). | 1c, 8c, 9c, 11c, 12c, 15c, 18c, 19c, 25c and 26c). |
The major axis PA (counted anticlockwise frou north to cast as usual) and € from the ellipse fitting are drawn asa function of the ellipse semiauajor axis (therefore. iu radius) iu Figs. | The major axis PA (counted anticlockwise from north to east as usual) and $\epsilon$ from the ellipse fitting are drawn as a function of the ellipse semi-major axis (therefore, in radius) in Figs. |
Le-28e. | 1e-28e. |
Iu Fies. | In Figs. |
1-28f the surface brightuess profiles Gsophotal magnitude fit to the ellipse ellipse semauajor axis) are given for I and J. We have used the 1-D resulting profiles to obtain the nlee and disk contributions. | 1f-28f the surface brightness profiles (isophotal magnitude fit to the ellipse ellipse semi-major axis) are given for K' and J. We have used the 1-D resulting profiles to obtain the bulge and disk contributions. |
We have fit an expoucutial to the outermost region (the disk). subtracted it to the observed profile aud fit art’! law (the bulge) to the residual: this calculation was continued uutil convergence was achieved (see Marquez Moles 1996. 1999). | We have fit an exponential to the outermost region (the disk), subtracted it to the observed profile and fit a $^{1/4}$ law (the bulge) to the residual; this calculation was continued until convergence was achieved (see Márrquez Moles 1996, 1999). |
The resulting bulge aud disk parameters are giveu -1 Table 3.. | The resulting bulge and disk parameters are given in Table \ref{decomposition}. |
The surface brightuess profiles in J aud I&. are dotted iu Fies. | The surface brightness profiles in J and K' are plotted in Figs. |
1£-28f with the bulec. disk aud bulee|disk fits superimposed. | 1f-28f with the bulge, disk and bulge+disk fits superimposed. |
The corresponding residuals are shown oei Fies. | The corresponding residuals are shown in Figs. |
Le-2Bse. Since residual backeround subtraction is 16 biggest source of error in determining the profiles (sec or iustauce de Jong 1996). we caution the reader that i1 je cases where the galaxy occupies most of the frame aud je residual backerounud level could only be determined iiu verv small regions. the error in the profiles can reach 15-20%. | 1g-28g. Since residual background subtraction is the biggest source of error in determining the profiles (see for instance de Jong 1996), we caution the reader that in the cases where the galaxy occupies most of the frame and the residual background level could only be determined in very small regions, the error in the profiles can reach . |
.. Otherwise. residuals are in general su:1ler thaw ©. except in the bar aud spiral aru reeions. | Otherwise, residuals are in general smaller than , except in the bar and spiral arm regions. |
rwermal relaxation phase. and that by the main sequence. all 1¢ simulations are almost indistinguishable. | thermal relaxation phase, and that by the main sequence, all the simulations are almost indistinguishable. |
Phe exception is the lowest. resolution simulation. | The exception is the lowest resolution simulation. |
The solid black line »egins its main sequence quite close to the dotted line (30 YO particle simulation). but has a hook near the turnolf. --=idicating the presence of a central convection zone. | The solid black line begins its main sequence quite close to the dotted line (30 000 particle simulation), but has a hook near the turnoff, indicating the presence of a central convection zone. |
As 10 resolution increases. we see that this convection zone isappears. | As the resolution increases, we see that this convection zone disappears. |
η SimulationsTη with. 10L particles. are not quiteη etailed enough to accurately depict the true structure of rese collision products. | Simulations with $10^4$ particles are not quite detailed enough to accurately depict the true structure of these collision products. |
The criterion for convective stability in stellar models is the Schwarzschild criterion: where V. is the acliabatic temperature eracient in the star. ancl V is the radiative temperature gradient: where c is the Stefan Boltzmann constant. € is the eravitational constant. anda.L.P.M and T are the opacity. Luminosity. pressure. enclosed mass and temperature at that position in the star. | The criterion for convective stability in stellar models is the Schwarzschild criterion: where $\nabla_{ad}$ is the adiabatic temperature gradient in the star, and $\nabla$ is the radiative temperature gradient: where $\sigma$ is the Stefan Boltzmann constant, $G$ is the gravitational constant, and $\kappa, L, P, M$ and $T$ are the opacity, luminosity, pressure, enclosed mass and temperature at that position in the star. |
The mocels taken from the SPLL results were used as starting models in the stellar evolution code. and allowed: to relax. so that all the equations of stellar structure are satisfied. | The models taken from the SPH results were used as starting models in the stellar evolution code, and allowed to relax, so that all the equations of stellar structure are satisfied. |
Since our SPLL code does not allow for energy transport. the luminosity distribution in the star can only be determined from the other structure parameters using the equations of stellar structure. | Since our SPH code does not allow for energy transport, the luminosity distribution in the star can only be determined from the other structure parameters using the equations of stellar structure. |
The temperature eracients were calculated. and their dilference is plotted as a [function of mass fraction in Figure 3.. | The temperature gradients were calculated, and their difference is plotted as a function of mass fraction in Figure \ref{convec}. |
The line stvles are the same as in Figure 1.. | The line styles are the same as in Figure \ref{structure}. |
We can clearly see a convective core in these stars. where YoWaar27 0. | We can clearly see a convective core in these stars, where $\nabla - \nabla_{ad} \ > 0$ . |
However. the entire star outside of Che inner 0.1 AZ. is not convective. | However, the entire star outside of the inner 0.1 $M_{\odot}$ is not convective. |
Phe outermostshell in all our stellar models is at a mass fraction Al/Al, 0.99667. | The outermostshell in all our stellar models is at a mass fraction $M/M_{\star} = 0.99667$ . |
This value is the same for all models. regardless of the number of SPLL particles. because of the way we have transformed the SPILL information to the stellar evolution starting model. | This value is the same for all models, regardless of the number of SPH particles, because of the way we have transformed the SPH information to the stellar evolution starting model. |
Phe total mass of these stars is 1.12 M... so only the outer 0.004 AL. is not resolved. | The total mass of these stars is 1.12 $M_{\odot}$, so only the outer 0.004 $M_{\odot}$ is not resolved. |
Therefore. we can sav with certainty that this stellar collision product does not have à surface convectionzone larger than 0.004 Αι. | Therefore, we can say with certainty that this stellar collision product does not have a surface convectionzone larger than 0.004 $M_{\odot}$ . |
Our | Our |
Within the hierarchical galaxy formation model. Dark Matter. (hereafter DAL) haloes are thought to play. the role of gravitational buiding blocks. within which barvonic dilfuse matter collapses :uxd becomes detectable. | Within the hierarchical galaxy formation model, Dark Matter (hereafter DM) haloes are thought to play the role of gravitational building blocks, within which baryonic diffuse matter collapses and becomes detectable. |
On ealactic scales. the formation of stars and their evolution provides an important. probe of he evolution of the visible content of the Universe (2?).. :though the subtleties of the stellar formation processes wihin galaxies. as of today not vet completely understood. hinders an exploitation of these objects as a clean probe of the evolution of DAL haloes. | On galactic scales, the formation of stars and their evolution provides an important probe of the evolution of the visible content of the Universe , although the subtleties of the stellar formation processes within galaxies, as of today not yet completely understood, hinders an exploitation of these objects as a clean probe of the evolution of DM haloes. |
On the other extreme of the mass scale. the most massive clusters of galaxies are regarded as one of the most reliable cosmological probes (?):: in. particular. their abundance and evolution with redshift is a very sensitive test of the underlying cosmological model The Mass Function (hereafter ME) is an indirect test of the total virialised mass of DAL haloes: exact. predictions of the latter can be done using the nonlinear spherical collapse model (?).. an essential ingredient. of the Press-Schechter | On the other extreme of the mass scale, the most massive clusters of galaxies are regarded as one of the most reliable cosmological probes : in particular, their abundance and evolution with redshift is a very sensitive test of the underlying cosmological model The Mass Function (hereafter MF) is an indirect test of the total virialised mass of DM haloes: exact predictions of the latter can be done using the nonlinear spherical collapse model , an essential ingredient of the Press-Schechter |
is however not the case for XPE, mainly due to the fact that the wind rates can be two orders of magnitude higher than the EUV-driven rates, meaning that at the time of gap opening the mass of the draining inner disc and the accretion rate onto the star of the inner disc material remain detectable for a non-negligible amount of time. | is however not the case for XPE, mainly due to the fact that the wind rates can be two orders of magnitude higher than the EUV-driven rates, meaning that at the time of gap opening the mass of the draining inner disc and the accretion rate onto the star of the inner disc material remain detectable for a non-negligible amount of time. |
For the disc population generated in this work, Figure 12,, shows that the accreting inner holes and non-accreting inner-holes (M <1x10Mo !)) are in general equally likely out to a radius of Rin~ 5AU. | For the disc population generated in this work, Figure \ref{fig:prob_holes}, shows that the accreting inner holes and non-accreting inner-holes $\dot{M}$ $<1\times10^{-11}$ ) are in general equally likely out to a radius of $R_{in}\sim5$ AU. |
Clearly as the transition disc is further photoevaporated and its inner radius moves out the accretion signatures onto the star become less evident and non-accreting inner holes dominate at radii larger than 20 AU. | Clearly as the transition disc is further photoevaporated and its inner radius moves out the accretion signatures onto the star become less evident and non-accreting inner holes dominate at radii larger than 20 AU. |
The total integrated ratio out to 10AU (the radius probed by 244m emission around solar-type stars) is found to be 2596 accreting and 7596 non-accreting for the entire population. | The total integrated ratio out to 10AU (the radius probed by $\mu$ m emission around solar-type stars) is found to be $25\%$ accreting and $75\%$ non-accreting for the entire population. |
We caution that this isnot equivalent to the observed fraction of accreting to non-accreting objects in a individual cluster, where the cluster age should also be accounted for. | We caution that this is equivalent to the observed fraction of accreting to non-accreting objects in a individual cluster, where the cluster age should also be accounted for. |
In young clusters the transition disc population is dominated by high X-ray luminosity objects which give rise to a considerably longer accreting inner hole phase. | In young clusters the transition disc population is dominated by high X-ray luminosity objects which give rise to a considerably longer accreting inner hole phase. |
In contrast this ratio is much lower in old clusters where the transition disc population is dominated by low ray luminosity objects that have very short accreting inner-hole phases. | In contrast this ratio is much lower in old clusters where the transition disc population is dominated by low X-ray luminosity objects that have very short accreting inner-hole phases. |
It is perhaps worth noting at this point that the detection of a ‘transition’ disc observationally is made through observation of the dust continuum spectral energy distribution (SED). | It is perhaps worth noting at this point that the detection of a `transition' disc observationally is made through observation of the dust continuum spectral energy distribution (SED). |
Alexander Armitage (2007) examined the behaviour of dust in a photoevaporating disc, finding that, under the action of dust drag, the time-scale for dust grains to drain onto the star is of order 10° yrs, after the gap opens, approximately two order of magnitudes faster than the gas draining time-scales. | Alexander Armitage (2007) examined the behaviour of dust in a photoevaporating disc, finding that, under the action of dust drag, the time-scale for dust grains to drain onto the star is of order $^3$ yrs, after the gap opens, approximately two order of magnitudes faster than the gas draining time-scales. |
This means that an observer would certainly see a significant drop in opacity in the inner disc immediately after a gap has opened, while the gas will still linger in the inner dust disc for the duration of its viscous draining time-scale of ~10? yrs. | This means that an observer would certainly see a significant drop in opacity in the inner disc immediately after a gap has opened, while the gas will still linger in the inner dust disc for the duration of its viscous draining time-scale of $\sim10^5$ yrs. |
We have used our population synthesis model to investigate the accretion rate versus inner hole size evolution for transition discs created by XPE under the assumption of immediate dust clearing at the time of gap opening. | We have used our population synthesis model to investigate the accretion rate versus inner hole size evolution for transition discs created by XPE under the assumption of immediate dust clearing at the time of gap opening. |
Figure 13 shows the probability distribution of the disc | Figure \ref{fig:innerholes} shows the probability distribution of the disc |
A limitation of the model in the stellar case is the ow power that can be couveved outo the chromosphere. estimated of the order of 1010Lo" NV, The observed phase lag between the plauet and the hot spot is also dificult to explain with a potential dipole maguctic field ike that of Jupiter. | A limitation of the model in the stellar case is the low power that can be conveyed onto the chromosphere, estimated of the order of $10^{16}-10^{17}$ W. The observed phase lag between the planet and the hot spot is also difficult to explain with a potential dipole magnetic field like that of Jupiter. |
MeIvoretal.(2006) sugeested hat the axis of the stellar dipole is tilted with respect o the orbital angular momentum of the planet. but even with this livpotliesis it is uot possible to accouut Or a phase lag of ~160? as observed in the case of e Aud. | \citet{McIvoretal06} suggested that the axis of the stellar dipole is tilted with respect to the orbital angular momentum of the planet, but even with this hypothesis it is not possible to account for a phase lag of $\sim 160^{\circ}$ as observed in the case of $\upsilon$ And. |
On the other haud. the model by Preusseetal— (2006). recently coufinued by the iuunerical sinulatious of oppetal.(2011)... can explain a large plase lag. | On the other hand, the model by \citet{Preusseetal06}, recently confirmed by the numerical simulations of \citet{Koppetal11}, can explain a large phase lag. |
It is based on packets of Alfven waves that are excited by the orbital motion of the planct. even iu the absence of an intrinsic planetary field. and then propagate towards the star along characteristics making some angle with the magueticfield lines. | It is based on packets of Alfven waves that are excited by the orbital motion of the planet, even in the absence of an intrinsic planetary field, and then propagate towards the star along characteristics making some angle with the magnetic field lines. |
They can zeach— the surface provided that the velocityof the stellar wind at the distance of the planet is subalfvenic. that is the case when the planet is sufficiently. close to the star. | They can reach the surface provided that the velocity of the stellar wind at the distance of the planet is subalfvenic, that is the case when the planet is sufficiently close to the star. |
The main limitation is the low cnerey fux at the stellar surface that males almost impossible to account for an emitted⋅ power as hiel. as QUO10291023oy WLXu To explain the emitted power. Lanza—(2009) considered the possibility that the interaction between the plauctary maeuctic ⋅⋅field aud the stellar coronal field may frigserao the release of the enerev already stored iu the coronal loops. | The main limitation is the low energy flux at the stellar surface that makes almost impossible to account for an emitted power as high as $10^{20}-10^{21}$ W. To explain the emitted power, \citet{Lanza09} considered the possibility that the interaction between the planetary magnetic field and the stellar coronal field may trigger the release of the energy already stored in the coronal loops. |
The reconnection process continuously: produced. by the planetary field⋅ iu⋅ the outer corona tends to reduce the magnetic↜∙ enerev. oο 16 coronal field. | The reconnection process continuously produced by the planetary field in the outer corona tends to reduce the magnetic energy of the coronal field. |
⋅ SinceP the minimE enerev state for: a eiven total maguetic helicity is a linear force-frec ⋅ : ⋅ ⋅ ↴↸∖≼↧∙↑∐↸∖∶↴∙⊾↸∖∪↕⊔↸∖⊓⋅⋅↖↽∪↕↑∐↸∖∏↸∖∐↕∐∐∖↴∖↴↸⊳⋜⋯↴⋝↸∖≺∐∖↴∖↴↸⊳↥⋅↕↴⋈∖≼ ⋅ analyticallym DMeiveu the boundaryu coucitious⋅∙↴ at the stellar↴⋅ surface. | Since the minimum energy state for a given total magnetic helicity is a linear force-free field, the geometry of the field lines can be described analytically given the boundary conditions at the stellar surface. |
m An interesting ∙↴⋅↜∙property of ∙⊽≽⋅linear ∙⊳⋅⊳⋅force-free ⋅ ↴↸∖≺↧↴∖↴↕↴∖↴↑∐⋜↧↑∪⋟↴⋝↸∖↕↕∶↴∙⊾↑↖↖↽↕↴∖↴↑↸∖≼↧∙↑∐∏↴∖↴↻↥⋅∪↖⇁↕≼∐∐∶↴⋁⋜↧∐⋜↧⊓∐⋅⋜↧ explanation for the phase lag between the planet and the chromospheric hot. spot located close to the ootpoiuts of the magnetic field lines. | An interesting property of linear force-free fields is that of being twisted, thus providing a natural explanation for the phase lag between the planet and the chromospheric hot spot located close to the footpoints of the magnetic field lines. |
This property ji been explored iu etail by Lanza(2008). | This property has been explored in detail by \citet{Lanza08}. |
. Iu conibination with the trigecring of magnetic energv release sugeested by Lanza—(2009).. force-free fields could account for the enütted power of the hot spots as well as for some coronal cussion culamccment. | In combination with the triggering of magnetic energy release suggested by \citet{Lanza09}, force-free fields could account for the emitted power of the hot spots as well as for some coronal emission enhancement. |
Depending ou the total magnetic helicity and the ⋅⋅ ⋅⋅ ⋅ ≻↥⋅↸∖≼∐↸⊳↑↑↕∐∖↥⋅⋜∥∐∪↸∖∐∐↴∖↴↴∖↴↕∪∐↕↥⋅∪⋯↑∐↸∖↴∖↴⋅↖↽↴∖↴↑↸∖⋯↴∖↴∐∪↴∖↴⊓∐∶↴∙⊾ ↕∐↑↸∖↥⋅↴∖↴↸∖↸⊳↑↕∐∶↴⋁↑↕∐∖↻↕⋜⋯↸∖↑⋜∐⋅⋅↖⇁⋯⋜↧∶owetosphere mav close up before. reaching. the surface. of. the star whichM may inplv that no magnetic connection exists between the planct and the chromosphere (seeLanza2009). | Depending on the total magnetic helicity and the surface boundary conditions, the magnetic field lines intersecting the planetary magnetosphere may close up before reaching the surface of the star which may imply that no magnetic connection exists between the planet and the chromosphere \citep[see ][]{Lanza09}. |
. Such configurations would eive rise to an "off state of the SPAT in the terminology of Shkoluikctal.(2008). | Such configurations would give rise to an "off" state of the SPMI in the terminology of \citet{Shkolniketal08}. |
When the boundary couditious and/or the total helicity change. a field configuration that counects the planet with the star can be restumed leading to a switcling on of the SPMI features in the chromosphere. | When the boundary conditions and/or the total helicity change, a field configuration that connects the planet with the star can be resumed leading to a switching on of the SPMI features in the chromosphere. |
Auother interesting property of linear force-free fields is the xedonminauce of closed field lines that reduce the loss of angular inonieutuni in the stella wind aud the cousequeut braliug of the stellar rotation. | Another interesting property of linear force-free fields is the predominance of closed field lines that reduce the loss of angular momentum in the stellar wind and the consequent braking of the stellar rotation. |
This may lave miportant consequences for the estimate of the age of the stars accompanied by lot Jupiters through he method of gvrochrouology (Lanza2010).. | This may have important consequences for the estimate of the age of the stars accompanied by hot Jupiters through the method of gyrochronology \citep{Lanza10}. . |
Full maguetohydrodyuanuc snaulatious. such a5 hose of Ipetal.(2001).. Cohenetal.(20002).. Cohenοἳal.(2009h).. Cohenetal.(2010).. Cohenctal.2011)... and Cohenetal(20115)... confirms several results of the analytical force-free models by Lanza(2008.2009.2010) and show how a stellar corona tends to be confined aud heated by the interaction with the Wagnetosphere ofa hot Jupiter. | Full magnetohydrodynamic simulations, such as those of \citet{Ipetal04}, \citet{Cohenetal09a}, \citet{Cohenetal09b}, \citet{Cohenetal10}, \citet{Cohenetal11a}, and \citet{Cohenetal11b}, confirm several results of the analytical force-free models by \citet{Lanza08,Lanza09,Lanza10} and show how a stellar corona tends to be confined and heated by the interaction with the magnetosphere of a hot Jupiter. |
Hot coronal regions can 2ccouut for an increased enission in the N-rays aud the ↥⋅↸∖≼⇂⋯⊳↸∖≼↧↸∖↕−⊔↸⊳↕↸∖↕⊔⊳⋅↖⇁∪↕⋟∏∐∖⋜↧↸⊳↸⊳↸∖↕↸∖↥⋅⋜↧↕∪∐∪≯↑↕∐∖↴ | Hot coronal regions can account for an increased emission in the X-rays and the reduced efficiency of the acceleration of the stellar wind. |
∖↴↑↸∖∐⋜∐⋅↖↖↽↕↓≼↧∙ Although the maguctolhvdrodyuamic regimes accessible ο nunierical simulations are many orders of magnitude ar from real systems. these studies provide a wealth of information ou what is to be expected in real star-lancet interaction aud cau be specialized to the case of articnlar ∖∖systems (ef.es...Cohen∢∖∖⋠al.≮≻↽↽⊀30111) | Although the magnetohydrodynamic regimes accessible to numerical simulations are many orders of magnitude far from real systems, these studies provide a wealth of information on what is to be expected in real star-planet interaction and can be specialized to the case of particular systems \citep[cf., e.g., ][]{Cohenetal11a}. |
Another approach to the star-planetet interaction lias considered individual coronal loops reconnecting with | planetary ↸∖⋪↸∖⋆⋅↽⋪∩⊾↸∖∪↴∖↴↸∖⋅↸∖∙⇀↔∪⋅≺∎∩⊾∪↸∖⋆↴∖↴↸∖magnetosphere, | Another approach to the star-planet interaction has considered individual coronal loops reconnecting with the planetary magnetosphere. |
According t th pha of the stellar osactivity evele. the height.o of the tops] of the loops] can‘ varyUM leadingBN im some cases‘ to ‘a strong interaction. modulated with the orbital period ral pe1 ∪↑∐∖≻⋪⋯ | According to the phase of the stellar activity cycle, the height of the tops of the loops can vary leading in some cases to a strong interaction, modulated with the orbital period of the planet. |
↸∖↑∙∪∐↑∐∖∪↑∐∖↥⋅↕⋪⋯≺∙∏∐⋅∐↕∩⊾⊔⋪↧↴∖↴↸∖↴∖↴↖↖↽∐∖∐ »H her land. duie iPASC 1 i6 loops are uot so tall the osinteraction with the ≻↕⋜⋯↸∖⋜∐⋅↖⇁∱∐∖↕≼↧↕↴∖↴↥⋅↸∖≼↧⋯⊳↸∖≼↧⋜⋯≼↧↖↖↽↸∖↻↥⋅↸∖≼↧∪∐∐∐⋜⊔↑↕↖↽↴∖↴↸∖↸∖ . ∐∖⋯∪≼⊔∏⋜↧↑↕∪∐↴⋝∙↖↽↑∐↸∖↴∖↴↑↸∖∐⋜∐⋅↥⋅∪↑⋜↧↑↕∪∐⋖≼↥⋅⋜⋯⋯↸∖↥⋅∙∖↽∺⋜↧⋜∐⋅ . )Ww 07). | On the other hand, during phases when the loops are not so tall the interaction with the planetary field is reduced and we predominantly see the modulation by the stellar rotation \citep{CranmerSaar07}. |
.. Other studies have addressed ie transportof energv from the recounection site. westunably iu the corona. to the chromosphere through vcauns of energetic particles that produce a localized ieatius when they inipact outo the chromosphere (Cu&Suzuki 2009). | Other studies have addressed the transportof energy from the reconnection site, presumably in the corona, to the chromosphere through beams of energetic particles that produce a localized heating when they impact onto the chromosphere \citep{GuSuzuki09}. |
. Finally. we should mention the inodoels developed to ⋅ . ⋅⋅ ↴∖↴↿∐⋅↕⋜⋯∖↴⋝∪∏∐≺↧⋜∐⋅⋅↖↽↸⊳∪∐≼∐⊓∪∐↴∖↴∙↑∐∖⋯⋜↧∶↴∙⊾∐↸∖↑↕↸⊳∏↸∖↕≼⊔∐∐∖↴∖↴⋅ . rot Jupiters. | Finally, we should mention the models developed to predict the radio emission from the systems hosting hot Jupiters. |
Since those planets are πιο closer than Jupiter. to the Sun..the stellar wind. speed is ⋅⋅Likely to )e 3ubalfveuic. | Since those planets are much closer than Jupiter to the Sun,the stellar wind speed is likely to be subalfvenic. |
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