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The sample of sources observed in this paper has been selected. among the brightest prototvpes [or the different twpes of nuclear activitv. | The sample of sources observed in this paper has been selected among the brightest prototypes for the different types of nuclear activity. |
. Their distances and corresponding linear scales are shown Table 1.. | Their distances and corresponding linear scales are shown Table \ref{tab:sourceandlines}. |
The sources can be grouped as follows. | The sources can be grouped as follows. |
AlS2 and253 are the strongest ancl richest extragalactic molecular sources. | and are the strongest and richest extragalactic molecular sources. |
These are the archetypes of starburst galaxies housing the two brightest extragalactic IRAS sources (Soilerοἱal.LO89). | These are the archetypes of starburst galaxies housing the two brightest extragalactic IRAS sources \citep{Soifer89}. |
. S82. as opposed to the voung starburst in 2253. is claimed to be the prototvpe of evolved starburst. | 82, as opposed to the young starburst in 253, is claimed to be the prototype of evolved starburst. |
This is supported by its observed chemistry characterized by low abundance of complex molecues such as SiO. ΟΙ (Alanersberger&IIenkelMartinetal.2006b) and large abundanceof molecular ions like . . and Πιο which are expected to be enhanced in PDRs (Fuente 2008). | This is supported by its observed chemistry characterized by low abundance of complex molecules such as SiO, $_3$ OH \citep{Mauers93,Martin06b} and large abundanceof molecular ions like $^+$, $^+$, and $_3$ $^+$ which are expected to be enhanced in PDRs \citep{Fuente05,Fuente06,vdtak08}. |
. In this work we did not observe the central position towards 882. but a position in (he north-east molecular complex (hereafter 582*) where most of the photodissociation regions are located as observed in the HCO emission high-resolution maps 2002)..2. | In this work we did not observe the central position towards 82, but a position in the north-east molecular complex (hereafter $^*$ ) where most of the photodissociation regions are located as observed in the HCO emission high-resolution maps \citep{Burillo02}. , |
is another wel studied starburst spiral galaxy. with a hieh molecular eas concentration in its nuclear region and showing traces of a tidal interaction with a cdwarl companion galaxy (Hurtetal.1996:Mason&Wilson2004). | is another well studied starburst spiral galaxy with a high molecular gas concentration in its nuclear region and showing traces of a tidal interaction with a dwarf companion galaxy \citep{Hurt96,Mason04}. |
..$2 is a nearby [aceou barred ealaxyv undergoing strong nuclear starburst likely due to gas inflow along its bar (Petitpas&al. 1979). | is a nearby faceon barred galaxy undergoing strong nuclear starburst likely due to gas inflow along its bar \citep{Petipas98,Talbot79}. |
. shows a moderate starburst in its nucleus also Claimed to be related to the presence of a bar (Turner&Ilo1983:Schinnereretal.2007). | shows a moderate starburst in its nucleus also claimed to be related to the presence of a bar \citep{Turner83,Schinne07}. |
. is a nearby luminous iirared galaxy. (LIRG) prototvpe of a Sevlert 2 nucleus (Antonueci&Miller1985). | is a nearby luminous infrared galaxy (LIRG) prototype of a Seyfert 2 nucleus \citep{Anto85}. |
. The AGN is enclosed by a circummuclear starburst ring with 14” (~ Lkkpe) radius (Myers&Scoville987:Sehinnereretal.2000). | The AGN is enclosed by a circumnuclear starburst ring with $14''$ $\sim1$ kpc) radius \citep{Myers87,Schinne00}. |
. Thus. we have observed two positions in 11068. | Thus, we have observed two positions in 1068. |
One towards the central AGN and an offset position 0".—14" towards a peak of emission within 1je cireummnuclear ring. hereafter. 11068*. mostly tracing the starburst. 4 | One towards the central AGN and an offset position $0'',-14''$ towards a peak of emission within the circumnuclear ring, hereafter $^*$, mostly tracing the starburst. , |
943. bei1 one of the three brightest IILAS point sources. has been the target of some a detailed moecular study by Wangetal. (2004)... | being one of the three brightest IRAS point sources, has been the target of some a detailed molecular study by \citet{Wang04}. . |
This nearly spiral galaxy. harbor a heavily obseured Sevlert 2 nucleus (Draatzetal.1997:Maiolino alsosurrounded by a starburst ring (Marconietal. 2000)... | This nearly edge-on spiral galaxy harbor a heavily obscured Seyfert 2 nucleus \citep{Braatz97,Maio99} alsosurrounded by a starburst ring \citep{Marconi00}. . |
551) is a | 51) is a |
The extrasolar planets discovered so far have masses. semimajor axes and eccentricities in the range 0.16.11 Jupiter masses (A44). 0.038447 au and Q00.93. respectively. | The extrasolar planets discovered so far have masses, semi--major axes and eccentricities in the range 0.16–11 Jupiter masses $M_{\rm J}$ ), 0.038–4.47 au and 0–0.93, respectively. |
Ln general. high. eccentricities are very. dillicult to explain in the context of a model where an isolated planet. forms in a disc. as the diseprotoplanet interaction leads to strong eccentricity damping (Nelson et al. | In general, high eccentricities are very difficult to explain in the context of a model where an isolated planet forms in a disc, as the disc–protoplanet interaction leads to strong eccentricity damping (Nelson et al. |
2000). | 2000). |
Papaloizouape ‘Terquem (2001.‘=⊳ hereaftert PTOL) investigated. a scenario D.in which. a population. of planetary mass objects.⊲ formed.⋅ rapidly. through a fragmentation⋅⊀ process occuring ⊀⊀in a thick. disc. or protostellar envelope on o scale of 100 au. | Papaloizou Terquem (2001, hereafter PT01) investigated a scenario in which a population of planetary mass objects formed rapidly through a fragmentation process occuring in a thick disc or protostellar envelope on a scale of 100 au. |
Such à svstem then underwent dynamical relaxation on a timescale of hundreds of orbits which resulted.. in ∢⊲⊀ejection of⋅ most of⋅ the objects. | Such a system then underwent dynamical relaxation on a timescale of hundreds of orbits which resulted in ejection of most of the objects. |
⊲ Le was found⋅ that the characteristics∙nm of⋅ massive∢⊽ eccentric∙⊀ extrasolaruu. planets and the massive. MDhot. Jupiter"n observed.u so. far⋅∙ ⊔↓⊀↓⋏∙≟↓↥∣⋈⋅⋯∙≼∙∪⊔⊔⊓⊾∠⊔⋅∪↓⋅∣⋡∙∖⇁⊳∖⋯⇍↓↕⋜ | It was found that the characteristics of massive eccentric extrasolar planets and the massive 'hot Jupiter' observed so far might be accounted for by such a model. |
↧⊔↓⋯⇂∢⊾↓⊳∐∪∖∖⋎⋖⋅∖⇁∢⊾↓⋅⊳↓≻↓⋜⋯∢⋅⇂⊳∖ with masses smaller than a few Ady are probably too small to have formed through a gravitational instability or fragmentation process (Masunaga Lnutsuka 1999. Boss 2000) and are thus more likely to have grown through the core accretion mechanism in à protoplanetary disc (Mizuno 1980). | However, planets with masses smaller than a few $M_{\rm J}$ are probably too small to have formed through a gravitational instability or fragmentation process (Masunaga Inutsuka 1999, Boss 2000) and are thus more likely to have grown through the core accretion mechanism in a protoplanetary disc (Mizuno 1980). |
llere. we investigate a scenario in which we have a planet accumulated. in a disc (PAID) on a timescale of 10" vears. being the observed lifetime of protostellar discs (Llaisch. Lada Lada 2001). together with a populationof massive outer planets already formed through fragmentation on a processesmuch shorter timescale. | Here, we investigate a scenario in which we have a planet accumulated in a disc (PAID) on a timescale of $10^6$ years, being the observed lifetime of protostellar discs (Haisch, Lada Lada 2001), together with a population of massive outer planets already formed through fragmentation processes on a much shorter timescale. |
During its formation phase. the PAID. is kept in a circular orbit by tidal interaction with the disc while the svstem of outer planets undergoes dynamical relaxation. | During its formation phase, the PAID is kept in a circular orbit by tidal interaction with the disc while the system of outer planets undergoes dynamical relaxation. |
After disc dispersal occurring at |=10" vears. the eecentricity of the PAID can be pumped up to high. values bv interaction. with. the remaining.. bound massive⋠⋎ planets. | After disc dispersal occurring at $t=10^6$ years, the eccentricity of the PAID can be pumped up to high values by interaction with the remaining bound massive planets. |
In order to focus∙ on the interaction| mechanisms.| we analvse the motion. of: a planet under a distant. perturber in. aE highly"n eccentric6006000dv orbit“hyd in1 wD$2. | In order to focus on the interaction mechanisms, we analyse the motion of a planet under a distant perturber in a highly eccentric orbit in \ref{sec:secular}. |
In So8/3. weuq presentDETES numericalyee simulations of the evolution of à svstenr containing outer massive planets and an inner PALD interacting with each ∪↥↓↥∢⋅ | In \ref{sec:results} we present numerical simulations of the evolution of a system containing outer massive planets and an inner PAID interacting with each other. |
↓⋅⋡↓⊲⇝↓⊔⋜↧∐∙∖⇁⊳⊲↓⊔≿∎≟∖∖⊽⋖⊾∠∐⊳∖≼∼⊔⊳∖⊳∖⋯⊔⋅↓⋅∢⊾⊳∖⊔↓↿⊳∖ | Finally, in \ref{sec:disc} we discuss our results. |
↓⋅⋡↓⊲⇝↓⊔⋜↧∐∙∖⇁⊳⊲↓⊔≿∎≟∖∖⊽⋖⊾∠∐⊳∖≼∼⊔⊳∖⊳∖⋯⊔⋅↓⋅∢⊾⊳∖⊔↓↿⊳∖⊳ | Finally, in \ref{sec:disc} we discuss our results. |
system, then the V magnitude and B—V colour outside eclipse would indeed be consistent with a distance near kkpc. | system, then the $V$ magnitude and $B-V$ colour outside eclipse would indeed be consistent with a distance near kpc. |
The yellow-red spectra of 2223 display a rather strong 455412 feature, as well as a weak (EW ~0.4 AA), but definitely detected 4 55592 line (see refspec223)). | The yellow-red spectra of 223 display a rather strong $\lambda$ 5412 feature, as well as a weak (EW $\sim 0.4$ ), but definitely detected $\lambda$ 5592 line (see \\ref{spec223}) ). |
Whilst there are no known interstellar features affecting the line, there are three, rather weak, DIBs at AA55404.5, 5414.8, and 5420.2 that are blended with 125412. | Whilst there are no known interstellar features affecting the line, there are three, rather weak, DIBs at $\lambda\lambda$ 5404.5, 5414.8, and 5420.2 that are blended with $\lambda$ 5412. |
To assess the real strength of the 4 55412 feature, we first evaluated the ratio of the EWs for some nearby DIBs (AA 55450.3, 5487.5) in the spectrum of MSP2223 and the corresponding EWs in the spectrum of 1183143 quoted by Herbig (1995)). | To assess the real strength of the $\lambda$ 5412 feature, we first evaluated the ratio of the EWs for some nearby DIBs $\lambda\lambda$ 5450.3, 5487.5) in the spectrum of 223 and the corresponding EWs in the spectrum of 183143 quoted by Herbig \cite{Herbig}) ). |
This ratio (1.33) was then used as a linear scaling factor to estimate the contamination of the 4 55412 line bythe three DIBs. | This ratio (1.33) was then used as a linear scaling factor to estimate the contamination of the $\lambda$ 5412 line bythe three DIBs. |
The corrected EW of the line amounts toAA. | The corrected EW of the line amounts to. |
. Therefore, 455412 is of comparable strength to the 255876 line (EW = ΑΑ)), although the latter is slightly affected (probably less than AA)) by a nebular emission line. | Therefore, $\lambda$ 5412 is of comparable strength to the $\lambda$ 5876 line (EW = ), although the latter is slightly affected (probably less than ) by a nebular emission line. |
These properties suggest an O7-8 spectral type (Walborn 1980)) for the combined spectrum. | These properties suggest an O7-8 spectral type (Walborn \cite{NRW}) ) for the combined spectrum. |
That we are dealing with an O-type star is confirmed by the absence of 144921 and the fact that the 2 66678 line is clearly blended with 2 66683. | That we are dealing with an O-type star is confirmed by the absence of $\lambda$ 4921 and the fact that the $\lambda$ 6678 line is clearly blended with $\lambda$ 6683. |
We measured the RVs of several stellar lines. | We measured the RVs of several stellar lines. |
The most robust results were obtained for 4 55412 and Ha, although we caution that the former is affected by blends with DIBs and for the latter, we had to perform a two Gaussian fit to account for the contamination by the nebular Ha emissionline®. | The most robust results were obtained for $\lambda$ 5412 and $\alpha$, although we caution that the former is affected by blends with DIBs and for the latter, we had to perform a two Gaussian fit to account for the contamination by the nebular $\alpha$ emission. |
. The results are listed in reftablemsp223 and plotted in refRVs223,, after subtracting the mean RV individually for each line. | The results are listed in \\ref{tablemsp223} and plotted in \\ref{RVs223}, , after subtracting the mean RV individually for each line. |
No obvious trend is apparent in this figure. | No obvious trend is apparent in this figure. |
The rather large shift in RV between the 455412 and Ha line is most likely due to wind emission affecting the latter line. | The rather large shift in RV between the $\lambda$ 5412 and $\alpha$ line is most likely due to wind emission affecting the latter line. |
Similar shifts in the apparent systemic velocity of different lines are commonly found in binary systems where at least one component features a strong stellar wind RRauw et citehde228766)). | Similar shifts in the apparent systemic velocity of different lines are commonly found in binary systems where at least one component features a strong stellar wind Rauw et \\cite{hde228766}) ). |
Our previous photometric monitoring of the cluster indicated that 2223 is probably an eclipsing binary (Paper I). | Our previous photometric monitoring of the cluster indicated that 223 is probably an eclipsing binary (Paper I). |
The Fourier analysis yielded the highest peak for a period of ddays, with a rather large uncertainty of ddays. | The Fourier analysis yielded the highest peak for a period of days, with a rather large uncertainty of days. |
The light curve folded with this period yields only a primary eclipse, with no indication of a secondary eclipse, although we might have missed it, if its duration is sufficiently short to fit into a small gap in phase coverage around ¢= 0.5. | The light curve folded with this period yields only a primary eclipse, with no indication of a secondary eclipse, although we might have missed it, if its duration is sufficiently short to fit into a small gap in phase coverage around $\phi = 0.5$ . |
In Paper I, we inferred an orbital periodof ddays, which then indicates two rather similar eclipses. | In Paper I, we inferred an orbital periodof days, which then indicates two rather similar eclipses. |
This would imply | This would imply |
This nuage had Όσοι obtained with the WEPC22 aboard he IIubble Space Telescope (IST). | This image had been obtained with the 2 aboard the Hubble Space Telescope (HST). |
It was taken with he FOOGW filter (Ape,À.. FWIIM-1578.7Aj). | It was taken with the F606W filter $\lambda_{mean}$, ). |
Further details are eiven in Diretta et al. (1996)). | Further details are given in Biretta et al. \cite{Biretta}) ). |
Additionally, further WEPC22 images of 11386 )ecanmie available from a survey of Sevtert ealaxics carried out by Wilson et al. | Additionally, further 2 images of 1386 became available from a survey of Seyfert galaxies carried out by Wilson et al. |
G6119). | 6419). |
These images aad been taken iu two narrowband aud two broadband filters. | These images had been taken in two narrowband and two broadband filters. |
In our studyv we nude use of the F517M aud F658N filter nuages. | In our study we made use of the F547M and F658N filter images. |
The nurowbaud (F658N) nuage covers the Ea r1] cmission within the filter FWA ofÀ. | The narrowband (F658N) image covers the $\alpha$ ] emission within the filter FWHM of. |
. The WFEPC22 images of L1386 were used to ideutifv the individual commponcuts of the NLR iu our 2-D spectrmm which will be described iu detail iu Sect. | The 2 images of 1386 were used to identify the individual components of the NLR in our 2-D spectrum which will be described in detail in Sect. |
1.3. | 4.3. |
Iu addition to the pure cussion line fluxes the calibrated spectrum also contains the radiation from the underlying stellar population (coutimmiun). | In addition to the pure emission line fluxes the calibrated spectrum also contains the radiation from the underlying stellar population (continuum). |
To remove this effect iu a siurple approach we caleulated a linear σωμα fit to our Spectruni across the oenüsson liues of interest Wa. uj). | To remove this effect in a simple approach we calculated a linear continuum fit to our spectrum across the emission lines of interest $\alpha$, ]). |
This linear interpolation which based on the mean flux in coutinmun windows (10 width) at either side of the Hà [Ni] doublet eiissiou line complex (AuAL. AU À)) was determined for each spatial scan aud subtracted. | This linear interpolation which based on the mean flux in continuum windows $\approx$ width) at either side of the $\alpha$ ] doublet emission line complex $\lambda_{blue}^c$, $\lambda_{red}^c$ ) was determined for each spatial scan and subtracted. |
The 2-D echelle spectrogram was transformed from wavelength space into velocity space as described by Dietrich Waencer (1998)). | The 2-D echelle spectrogram was transformed from wavelength space into velocity space as described by Dietrich Wagner \cite{DW}) ). |
We used the helioceutric velocity e,=918d31s+ of 111560, as giveu by IIuchtiieier Richter (1989)). | We used the heliocentric velocity $v\mathrm{_{sys}=918\pm34\,km\,s^{-1}}$ of 1386, as given by Huchtmeier Richter \cite{HR}) ). |
We checked this value by nieasuriue the ccutroid of the redshifted interstellar Π line and appliug the appropriate hehoceutric correction. | We checked this value by measuring the centroid of the redshifted interstellar I line and applying the appropriate heliocentric correction. |
This results iu ey,=921dS51kins+. in sood agreement with Uneltineier Richter (1989)). | This results in $v\mathrm{_{\rm sys}
= 924\pm51\,km\,s^{-1}}$, in good agreement with Huchtmeier Richter \cite{HR}) ). |
Since our spectra covers he waveleneth rauge1125À. we have several strong ciissiou lines which are useful for he kinematical study of the Lue cmitting gas. | Since our spectrum covers the wavelength range, we have several strong emission lines which are useful for the kinematical study of the line emitting gas. |
We use Πα aud the [Nu] lines for our. kinematical investigation. | We use $\alpha$ and the ] lines for our kinematical investigation. |
The less iuteuse AA66717.6731 doublet. A 66300. aud the AA 55721.6087. lines will be used for the investigation of the spatial variations of the physical couditious in the individual NLR cussionliue chumps. | The less intense $\lambda\lambda$ 6717,6731 doublet, $\lambda$ 6300 and the $\lambda\lambda$ 5721,6087 lines will be used for the investigation of the spatial variations of the physical conditions in the individual NLR emission–line clumps. |
By transforming the spectrum iuto the velocity space. the fe anap shown in Fie. 5 | By transforming the spectrum into the velocity space, the $l-v$ map shown in Fig. \ref{F5} |
ος has been obtained. | c has been obtained. |
At a distance of MMpe. the area covered by the slit of the spectrograph projected outo the plane of the galaxy is [80pe«60pc. | At a distance of Mpc, the area covered by the slit of the spectrograph projected onto the plane of the galaxy is $480\,\mathrm{pc} \times 60\,\mathrm{pc}$. |
The sinallest details visible in the 7—e map are estimated to be of the order of 0733 which corresponds to a linear size at the distance of 11386 of only ~Lsoc. | The smallest details visible in the $l-v$ –map are estimated to be of the order of 3 which corresponds to a linear size at the distance of 1386 of only $\sim 18\,\mathrm{pc}$. |
With this lnieh spatial resolution iu our spectra it Is possible to obtain a detailed picture of the kinematics of the NLR yomji LI386. based on a 2-D echelle spectzin with xitial information as a function of velocity. | With this high spatial resolution in our spectra it is possible to obtain a detailed picture of the kinematics of the NLR in 1386, based on a 2-D echelle spectrum with spatial information as a function of velocity. |
In Fig. | In Fig. |
3 the 4.oitial scans of the velocity space transformed spectrum LITEi the AG582. line are shown. | \ref{F3} the spatial scans of the velocity space transformed spectrum in the $\lambda$ 6583 line are shown. |
Iu order to nake a uautitative approach. a decomposition using elliptical Gaussians was applied to the spectrum. | In order to make a quantitative approach, a decomposition using elliptical Gaussians was applied to the spectrum. |
This vielded the position aud velocity of the individual componoeusin the fe inap as well as the line widths for cach iudividual cloud iu the |e space. | This yielded the position and velocity of the individual components in the $l-v$ –map as well as the line widths for each individual cloud in the $l-v$ space. |
The detailed description of the process of decomposing a 2-D echelle spectrogram is giveu bv Dietrich Waener (1998)) for the Sevtert ealaxy NGC110685. | The detailed description of the process of decomposing a 2-D echelle spectrogram is given by Dietrich Wagner \cite{DW}) ) for the Seyfert galaxy 1068. |
↽∕∏∐∖≼∐∖↸⊳∪∐∏⋯↴∖↴↕⊓∪∐∪↕⋟↑∐↸∖⋀∖⊽≼∶≼⊲∐∶≩≺∖∖⊓↴∖↴↻↸∖↸⊳⊓⋅⋜↧⋅↖↽↕↸∖↕≼∐∖≼↧ ∩↕∐≼∐↖↽↕≼↧∏⋜↧⋀∖⊽∫⇀↕⊰↸⊳↕∪⋯↧↴∖↴↖↖↽∐↕↸⊳∐↴∖↴∐∪↖↖⇁∐∐∖⋜⋯↖⇁↸∖↕∪↸⊳↕↑⋅↖⇁ ≼∐↕−↥∎↸∖↥⋅↸∖∐↸⊳↸∖↴∖↴↿⋃↑∪∐∣≼⊔↘↽⋯↴∖↴↓↕∐⊪ | The decomposition of the 1386 spectra yielded 9 individual NLR clouds which show mean velocity differences up to ${\rm{km\,s^{-1}}}$ in $\lambda$ 6583. |
∖⊽∐∫∕∖⊓⋅↱⊐≺∖∖∶≩∙⊺∐∖↕∐≼∐↖↽↕≼↧∏⋜↧↕ ⋀∖⊽∫⇀↕⊰↸⊳∪∐∏⋯∐∖∐↴∖↴∙↕⊔⋜∐⋅↨↘↽↸∖≼∏⋝∙↖⇁↕↸∖↑↑↸∖↥⋅↴∖↴∙⋜⋯≼↧≼∐∖↥⋅↕↖↽↸∖≺∏⋝∙↖↽↑∐↸∖ ≼∶⋜⋯↴∖↴↴ | The individual NLR components, marked by letters, and derived by the Gaussian decomposition (see Fig. \ref{F5} |
∖↴↕⋜⋯≼∐∖↸⊳∪∐∏⋯↴∖↴↕↑↕∪∐⋖↴∖↴↸∖↸∖∏∶↴∙⊾∙⋅↱⊐ ∙↸⊳↸⊳⋟⋜∐⋅↸∖↕∪↸⊳⋜↧↑↸∖≼↧⋜↧↑≼∐↕−↥∎↸∖↥⋅↸∖∐↑↖⇁↸∖↕∪↸⊳↕↑↕↸∖↴∖↴⋜⋯≺↧↴∖↴↻⋜↧↕⋜↧⊔∪↸⊳⋜↧↑↕∪∐↴∖↴ relative to the ↕⋯⊳↕↸∖∏↴ | c) are located at different velocities and spatial locations relative to the nucleus. |
∖↴∙↽∕∏∐∖∐∐↸∖↖↖↽↕≼⋯↓↴∖↴∪↕⋟↑↕∐∖↴∖↴↸∖↸⊳↕∪⋯↧↴∖↴ vary from abou ∩∪⊇∶≩⋅↱⊐↨↘↽⋯↴∖↴↓⋖↸⊳↕∪∏≺↧↴∖↴⊀≚∙↕≻⋅⇪≓↕⋝⋜⋯≼⇂ are accoluipane by clouds with considerably higher values of FWIIM = 270{7WOkms+ (clouds B. C. and E). | The line widths of these clouds vary from about $90-235\,{\rm{km\,s^{-1}}}$ (clouds A, D, F-I) and are accompanied by clouds with considerably higher values of FWHM = $270-470\,{\rm{km\,s^{-1}}}$ (clouds B, C, and E). |
The | The |
where (he subscript "ad refers to the adiabatic value. | where the subscript “ad” refers to the adiabatic value. |
Fig. | Fig. |
9. shows (he region. where (he mean superadiabaticdiiv peaks as a function of log P in the Sun and in Procvon. | \ref{sal} shows the region where the mean superadiabaticity peaks as a function of log P in the Sun and in Procyon. |
A major difference between the two simulations is that the peak of the SAL in Proevon is in a region of much lower density than in the Sun. | A major difference between the two simulations is that the peak of the SAL in Procyon is in a region of much lower density than in the Sun. |
The SAL peak in the Sun is located below the photosphere in a region that is relatively optically thick. and although the solar photosphere is affected by the convective flows just below. the position of the photospheric surface sullers little change due to the convective motions. | The SAL peak in the Sun is located below the photosphere in a region that is relatively optically thick, and although the solar photosphere is affected by the convective flows just below, the position of the photospheric surface suffers little change due to the convective motions. |
This is in sharp contrast with Procvon. where the more violent convective [flows are present in the atmosphere and cause ihe SAL position to vary with time. | This is in sharp contrast with Procyon, where the more violent convective flows are present in the atmosphere and cause the SAL position to vary with time. |
The SAL moves racially back and forth over a time of about. 20-30 minutes aud over a distance of up (to about half a local pressure scale heieht (about. 500 km). | The SAL moves radially back and forth over a time of about 20-30 minutes and over a distance of up to about half a local pressure scale height (about 500 km). |
By computing the Full Width at. Hall-Maxinum of the vertical velocity we obtain a mean granule diameter of about 10000km (details of how we estimate the granule diameter are given in Robinson et al. ( | By computing the Full Width at Half-Maximum of the vertical velocity we obtain a mean granule diameter of about 10000km (details of how we estimate the granule diameter are given in Robinson et al. ( |
2004). while the run of rms vertical velocity. Fig. | 2004), while the run of rms vertical velocity, Fig. |
5. provides a velocitv scale for the granules of 5-7 km/s. Combining these quantities vields an average lime scale. roughly half the time it takes the largest granules to overturn fully. in the vicinity ol 20-30 minutes. | \ref{vzrms} provides a velocity scale for the granules of 5-7 km/s. Combining these quantities yields an average time scale, roughly half the time it takes the largest granules to overturn fully, in the vicinity of 20-30 minutes. |
A 42 minute movie of granulation in Procvon for the box is presented at iibweastroaale.eduimarjf. | A 42 minute movie of granulation in Procyon for the box is presented at $ www.astro.yale.edu/marjf $. |
The largest granules (seen in the top right hand corner of the picture) last about 50 minutes. | The largest granules (seen in the top right hand corner of the picture) last about 50 minutes. |
The overturning granular motion causes the SAL peak to approach the optically thin lavers part of the time. | The overturning granular motion causes the SAL peak to approach the optically thin layers part of the time. |
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