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In order to obtain an atmosphere in radiative equilibrium. the energy transport part of the energy solver can be used instead of the temperature correction procedure. | In order to obtain an atmosphere in radiative equilibrium, the energy transport part of the energy solver can be used instead of the temperature correction procedure. |
The main problem ts that about a few hundred time steps are needed to obtain the resulting atmosphere structure in radiative equilibrium. while the temperature correction needs fewer Iteration steps and ts. therefore. significantly faster, | The main problem is that about a few hundred time steps are needed to obtain the resulting atmosphere structure in radiative equilibrium, while the temperature correction needs fewer iteration steps and is, therefore, significantly faster. |
(Thompsonetal.2007) 1.6/2 [να 1.67n ταν ! Matsumotoetal.(2005) ?7« tat Lym. Matsumotoetal.(2005) μι Εαν Salvaterra&Ferrara(2003))) | \citep{thm07} $\micron$ $\micron$ $\micron$ $^{-2}$ $^{-1}$ \citet{mat05} $^{-2}$ $^{-1}$ $\micron$ \citet{mat05} \citet{mat05} $\micron$ $\micron$ \citet{sal03}) |
ls ἰ1e so-called ion collisionless skin depth (ή2€>/m,c is the classical electron radius). | is the so-called ion collisionless skin depth $r_e\equiv e^2/m_e c^2$ is the classical electron radius). |
If thec«ndition (4.2)) is not satisfied. then the Sweet-Parker slow resistive reconnection is nol appicable and one may have a much faster collisionless reconnection regime (attributed to either Hall effect or anomolous resistivity). | If thecondition \ref{eq-condition-collisional}) ) is not satisfied, then the Sweet-Parker slow resistive reconnection is not applicable and one may have a much faster collisionless reconnection regime (attributed to either Hall effect or anomolous resistivity). |
Using equation (2.3)). we have | Using equation \ref{eq-delta_of_A}) ), we have. |
Thus. the condition that the laver is in the collisional regime. 9d;. can be cast as 7> 31(64) | Thus, the condition that the layer is in the collisional regime, $\delta> d_i$, can be cast as > )^4. |
Alore specificallv. in thebremsstrahlung case. using equation (2).1)) we see that the requirement (hat the laver is collisional reads "uL. Q4 boo μιGrfw(05) which is easily satisfied for many astroplivsical plasmas. | More specifically, in thebremsstrahlung case, using equation \ref{eq-brems-A}) ) we see that the requirement that the layer is collisional reads > _T^2 = _T^2, which is easily satisfied for many astrophysical plasmas. |
Sinularly. for externalinverse Compton cooling. using equation (3.2)). the collisional reconnection condition becomes (66 | Similarly, for externalinverse Compton cooling, using equation \ref{eq-EIC-A}) ), the collisional reconnection condition becomes > . |
Sinularly. for externalinverse Compton cooling. using equation (3.2)). the collisional reconnection condition becomes (66) | Similarly, for externalinverse Compton cooling, using equation \ref{eq-EIC-A}) ), the collisional reconnection condition becomes > . |
Based on the earlier studies of V1309. Ort (halter et al. | Based on the earlier studies of V1309 Ori (Shafter et al. |
1995: Schmidt Stockman 2001: Staude et al. | 1995; Schmidt Stockman 2001; Staude et al. |
2001) and our observations (83) it is clear that variations in the light curves are mostly caused by prominent stream emission with a smaller fraction due to exclotron emission. | 2001) and our observations 3) it is clear that variations in the light curves are mostly caused by prominent stream emission with a smaller fraction due to cyclotron emission. |
Lt is therefore necessary to include a large unpolarized background when modelling the polarised light curves. | It is therefore necessary to include a large unpolarized background when modelling the polarised light curves. |
We have done this in wo different wavs: (1) with constant background. as [irst approximation. and (2) adopting the observed total [ux as a [function of the orbital phase as the variable unpolarized xickeround. | We have done this in two different ways: (1) with constant background as first approximation, and (2) adopting the observed total flux as a function of the orbital phase as the variable unpolarized background. |
The latter case is more realistic for V1309 Ori due to the large amount of the unpolarized. [lux from he stream emission which varies during orbital evele. | The latter case is more realistic for V1309 Ori due to the large amount of the unpolarized flux from the stream emission which varies during orbital cycle. |
The diluting Hux from stream emission in V1309 Ori is 510 imes larger than the peak evelotron flux. estimated. [rom he fast drop of intensity at the eclipse of the compact source observed at phase &=0.952 (83.1). | The diluting flux from stream emission in V1309 Ori is 5–10 times larger than the peak cyclotron flux, estimated from the fast drop of intensity at the eclipse of the compact source observed at phase $\Phi=0.952$ 3.1). |
The X-ray temperature is chosen ki, LO keV according to observations of de Alartino et al. (1994).. | The X-ray temperature is chosen ${\rm kT_{brems}}$ = 10 keV according to observations of de Martino et al. \shortcite{b9}. |
Two-pole accretors. for example VV Pup (Wickramasinghe.Ferrario&Bailey 19890).. UZ For (Schwope.Beuermann& 1990).. DP Leo (Cropper&Wickramasinghe1993) and QS Tel (Schwopeetal.1995).. have been found to show dilferent. magnetic field. strengths in the accretion regions located in opposite hemispheres: up to factor of 2 dillerence in the field strengths has been measured. | Two-pole accretors, for example VV Pup \cite{b47}, UZ For \cite{b37}, DP Leo \cite{b7} and QS Tel \cite{b38}, have been found to show different magnetic field strengths in the accretion regions located in opposite hemispheres: up to factor of 2 difference in the field strengths has been measured. |
The more strongly accreting. pole has normally the weaker magnetic field. | The more strongly accreting pole has normally the weaker magnetic field. |
In the case of V1309. Ori there are no major cillerences seen in the wavelength. dependence of the positive and negative excursions of the observed: circular polarization curves (Figure 2)). suggesting that both accreting regions are accreting nearly equally. | In the case of V1309 Ori there are no major differences seen in the wavelength dependence of the positive and negative excursions of the observed circular polarization curves (Figure \ref{circpol}) ), suggesting that both accreting regions are accreting nearly equally. |
Pherefore. we have chosen the electron temperature and the plasma parameter (X —107) o be the same in both regions. | Therefore, we have chosen the electron temperature and the plasma parameter $\Lambda$ $^{5}$ ) to be the same in both regions. |
We assume an inclination of στο BStaude et al. (2001).. | We assume an inclination of $^{\circ}$ Staude et al. \shortcite{b42}. |
Parameters such as the ongitude of the emission region(s) on the surface of the white «να and. extension of the accretion regions were varied to match the gross. features. seen in the circular »olarization behaviour. | Parameters such as the longitude of the emission region(s) on the surface of the white dwarf, and extension of the accretion regions were varied to match the gross features seen in the circular polarization behaviour. |
Due to verv low level of the linear »olarization (less than 0.5 per cent). and noisy position angle variations. we have not tried to use linear polarization to ix anv model parameters. | Due to very low level of the linear polarization (less than 0.5 per cent), and noisy position angle variations, we have not tried to use linear polarization to fix any model parameters. |
Another reason for not. doing so is that scattering from free electrons of the stream can introduce significant linear polarization ellects. dominating over the low linearly polarized lux of evelotron origin. | Another reason for not doing so is that scattering from free electrons of the stream can introduce significant linear polarization effects, dominating over the low linearly polarized flux of cyclotron origin. |
Estimates of the colatitude of the accretion region «7 can be mace if sve assume that the observed circular polarization from different poles is not significantly alfected. by possible overlap of the polarized emission from. another accretion region. | Estimates of the colatitude of the accretion region $\beta$ can be made if we assume that the observed circular polarization from different poles is not significantly affected by possible overlap of the polarized emission from another accretion region. |
We can estimate for this region using the curation of the self-eclipse of the accretion region and equation (1) of Visvanathan Wickramasinghe (1981).. | We can estimate $\beta$ for this region using the duration of the self-eclipse of the accretion region and equation (1) of Visvanathan Wickramasinghe \shortcite{b44}. . |
Using the circular polarisation data shown in Figure 2.. we find of=145 [or the positive pole. and 3=35" for the negative pole. assuming an inclination 2—18 by Staucle et al. (2001).. | Using the circular polarisation data shown in Figure \ref{circpol}, we find ${\beta=145^\circ}$ for the positive pole, and ${\beta=35^\circ}$ for the negative pole, assuming an inclination ${\it i=\rm 78^{\circ}}$ by Staude et al. \shortcite{b42}. |
Observed circular polarization variations can be reproduced reasonably well with a model consisting of two separate emission regions. one centered at colatitude 3 = 145 (the xositive. pole). seen closest to the observer at @&=0.20. and another region centered in at 3 = 35 (the negative role}. seen closest to observer 9=0.70. | Observed circular polarization variations can be reproduced reasonably well with a model consisting of two separate emission regions, one centered at colatitude $\beta$ = $^{\circ}$ (the positive pole), seen closest to the observer at ${\Phi=0.20}$, and another region centered in at $\beta$ = $^{\circ}$ (the negative pole), seen closest to observer ${\Phi=0.70}$. |
For both regions we have adopted in Figure 10 longitudinal extension of 30° (in white dwarf rotational coordinates). but these values are not strongly. constrained. | For both regions we have adopted in Figure 10 longitudinal extension of $^{\circ}$ (in white dwarf rotational coordinates), but these values are not strongly constrained. |
Extensions in the range 101 60. eive almost similar results. | Extensions in the range $^{\circ}$ – $^{\circ}$ give almost similar results. |
Point-like emission region gives oo sharp polarization variations and very extended emission regions (larger than 607) too smooth and. low-amplitucde Curves. | Point-like emission region gives too sharp polarization variations and very extended emission regions (larger than $60^{\circ}$ ) too smooth and low-amplitude curves. |
The model shown in Figure 10. assumes that evelotron harmonics 6. 5. 4. 3. and 2. dominate in the CDVHU passbancs. respectively. | The model shown in Figure \ref{polmodel} assumes that cyclotron harmonics 6, 5, 4, 3, and 2, dominate in the $UBVRI$ passbands, respectively. |
This corresponds to a magnetic field of about 50 ALG. which is similar to the estimated values for magnetic field (33.55. MG. Garnavich et al. | This corresponds to a magnetic field of about 50 MG, which is similar to the estimated values for magnetic field (33– 55 MG Garnavich et al. |
1994: 61. M Shafter et al. | 1994; 61 MG Shafter et al. |
1995). | 1995). |
Fora 50 MC field the wavelengths of the xuwmonies 6 to 3 are at 3580A. 4300X. 5370A and 7160.A. i.c. one evelotron harmonic clearly dominates in cach of the (PVR bands. | For a 50 MG field the wavelengths of the harmonics 6 to 3 are at ${\rm 3580~\AA}$, ${\rm 4300~\AA}$, ${\rm 5370~\AA}$ and ${\rm 7160~\AA}$, i.e. one cyclotron harmonic clearly dominates in each of the $UBVR$ bands. |
Our 4 band. (8300 A) falls about half-way tween the 3rd. and 2nd harmonies at B~50ALG. | Our $I$ band ${\rm 8300~\AA}$ ) falls about half-way between the 3rd and 2nd harmonics at ${\rm B\sim50~MG}$. |
The vest correspondence to the observed. circular polarization variations is achieved using the model where unpolarized v~ackerouncl varies in à similar way to the total observed lux over the orbital evele (Figure LOL.continuous Linjg. | The best correspondence to the observed circular polarization variations is achieved using the model where unpolarized background varies in a similar way to the total observed flux over the orbital cycle (Figure \ref{polmodel},continuous line). |
Our model parameters for the location of accretion regions (in white chwarl rotation coordinates). positive vole αἱ 32145.V—70 and. negativo pole at j—-35.w-l110'. are similar to. the values. reported » Llarrop-Allin et al. (1999)... | Our model parameters for the location of accretion regions (in white dwarf rotation coordinates), positive pole at ${\rm \beta =
145^{\circ}, \Psi=-70^{\circ}}$ and negative pole at ${\rm \beta =
35^{\circ}, \Psi=110^{\circ}}$, are similar to the values reported by Harrop-Allin et al. \shortcite{b15}, |
who modelled white light data of Buckley Shafter (1995): 3=40° and 3=140 | who modelled white light data of Buckley Shafter \shortcite{b4}: $\beta = 40^{\circ}$ and $\beta = 140^{\circ}$. |
In contrast Staude et al. | In contrast Staude et al. |
(2001). derived: values of 3 = 17 and V=16" from their Doppler maps. | \shortcite{b42} derived values of $\beta$ = $^{\circ}$ and $\Psi=-16^{\circ}$ from their Doppler maps. |
In that studs it was assumed that only one accretion region was visible. | In that study it was assumed that only one accretion region was visible. |
Although it is possible for an accretion region to show both positive and negative circular polarisation. it is only for a short phase duration if we observe the ‘uncerside’ of the shock. | Although it is possible for an accretion region to show both positive and negative circular polarisation, it is only for a short phase duration if we observe the `underside' of the shock. |
ὃν comparing our circular polarization curves with those found in the literature we confirm that the spinof the white dwarl in V1309 Ori is svnchronised with the orbital period to a high degree. | By comparing our circular polarization curves with those found in the literature we confirm that the spinof the white dwarf in V1309 Ori is synchronised with the orbital period to a high degree. |
The zero-crossings of circular polarization take place at the same phase of the orbital period as found by Buckley Shafter (1995).. which puts an upper limit of | The zero-crossings of circular polarization take place at the same phase of the orbital period as found by Buckley Shafter \shortcite{b4}, which puts an upper limit of |
the ratio NOERIDZNGILEBLID tends to 0. | the ratio $N(-\mbox{FRII})/N(+\mbox{FRII})$ tends to $0$. |
We can define an 2. where Relativistic effects are niore important than intrinsic/environmental asvinmetries if QO<2xl1. while. Wool<20. the latter ellects are dominant. | We can define an $\varepsilon$, where Relativistic effects are more important than intrinsic/environmental asymmetries if $0<\varepsilon\leq1$, while, if $-1\leq\varepsilon<0$, the latter effects are dominant. |
Lf 0. the contribution of cach elfect is of comparable importance. | If $\varepsilon \sim 0$, the contribution of each effect is of comparable importance. |
Let us consider how the observed. distribution of η depends upon the properties of the asymmetric model. | Let us consider how the observed distribution of $x_{\rm l}$ depends upon the properties of the asymmetric model. |
The mean speeds of the jet and counterjet lobes are the sums of the mean intrinsic speed of the lobes. eo. and. the random dispersion about the mean. η=eo|cj and e=ος|ng. | The mean speeds of the jet and counterjet lobes are the sums of the mean intrinsic speed of the lobes, $v_0$ , and the random dispersion about the mean, $v_{\rm j} = v_0 + v_{\rm jd}$ and $v_{\rm cj} = v_0 + v_{\rm cjd}$. |
Η we assume that the intrinsic asymmetry. is distributed isotropically in the sky. the distribution function of the space velocities of the lobes on the jet- and counterjetsides should be the same. ο).=fey)Fin) implving that. the distribution functions Ge)=(eg)Gn). | If we assume that the intrinsic asymmetry is distributed isotropically in the sky, the distribution function of the space velocities of the lobes on the jet- and counterjet-sides should be the same, $F(v_{\rm j}) = F(v_{\rm cj}) = F(v_{\rm l})$, implying that the distribution functions $G(v_{\rm jd}) = G(v_{\rm cjd}) = G(v_{\rm d})$. |
For illustrative purposes. let us assume that (i) the distribution function of the intrinsic velocities ey is Gaussian with the mean Ty and standard deviation σι, and (ii) the distribution function. of random. disturbed: velocities is also à Gaussian with constant standard. deviation. σοι on the jet- and counterjet sides. | For illustrative purposes, let us assume that (i) the distribution function of the intrinsic velocities $v_0$ is Gaussian with the mean $\overline{v}_0$ and standard deviation $\sigma_{v_0}$ and (ii) the distribution function of random disturbed velocities is also a Gaussian with constant standard deviation $\sigma_{v_{\rm d}}$ on the jet- and counterjet sides. |
We assume that. these normal distribution functions are independent. | We assume that these normal distribution functions are independent. |
Then. the distribution function of their sum. the lobe speeds ο and Do are also normal with a mean speed Ty and. dispersion >2 For illustrative purposes. let us adopt à mean lobe speed vj=0.156 and a,=0.05c. | Then, the distribution function of their sum, the lobe speeds $v_{\rm j}$ and $v_{\rm cj}$ are also normal with a mean speed $\overline{v}_0$ and dispersion $\sigma_{v_{\rm l}}^2=\sigma_{v_0}^2+\sigma_{v_{\rm d}}^2$, For illustrative purposes, let us adopt a mean lobe speed $\overline{v}_{\rm l} = 0.15 c$ and $\sigma_{v_{\rm l}} = 0.05c$. |
Then. in Fig. | Then, in Fig. |
4 and Table 1. the predicted: distributions of wy. and the mean values of 2 and wy. are shown for different combinations of the intrinsic dispersions in the lobe speeds 7, and the intrinsic/environmental asymmetry 0,4. | 4 and Table 1, the predicted distributions of $x_{\rm l}$, and the mean values of $\varepsilon$ and $x_{\rm l}$, are shown for different combinations of the intrinsic dispersions in the lobe speeds $\sigma_{v_0}$ and the intrinsic/environmental asymmetry $\sigma_{v_{\rm d}}$ . |
The first and fourth lines of Table 1 correspond to the limiting cases of m,z0 ancl 0.05 respectively. | The first and fourth lines of Table 1 correspond to the limiting cases of $\sigma_{v_{\rm d}} \approx 0$ and 0.05 respectively. |
lt can be seen that the distribution function. of αι becomes more asvnimetrie for large values of 0,4. leading to a decrease in the value of mj. | It can be seen that the distribution function of $x_{\rm l}$ becomes more asymmetric for large values of $\sigma_{v_{\rm d}}$, leading to a decrease in the value of $\overline{x}_{\rm l}$. |
Therefore. the simple relation (6) between the mean speed. and the mean of the observed. fractional separation dillerence underestimates 7] and should be written The asymmetry. parameter is a better indicator of the intrinsic/environmental asvmnmetry than mj (see Table. 1). | Therefore the simple relation (6) between the mean speed and the mean of the observed fractional separation difference underestimates $\overline{v}_{\rm l}$ and should be written The asymmetry parameter is a better indicator of the intrinsic/environmental asymmetry than $\overline{x}_{\rm l}$ (see Table 1). |
Notice that the value 5zO is attained when σοι7044. that is. the intrinsic/environmental and relativistic cllects are of Comparable importance. | Notice that the value $\varepsilon \approx 0$ is attained when $\sigma_{v_{\rm d}} \approx \sigma_{v_0}$, that is, the intrinsic/environmental and relativistic effects are of comparable importance. |
lt should be noted.that the asvnimetry parameter eis a rather crude quantitative measure of the role of intvinsicfenvironmental asvmmetries because of thesmall number statistics generally. involved. in determining | It should be notedthat the asymmetry parameter $\varepsilon$ is a rather crude quantitative measure of the role of intrinsic/environmental asymmetries because of thesmall number statistics generally involved in determining |
whereas the ereen (light erev) curve is fit to the wid-to-late M dwarf suuple. | whereas the green (light grey) curve is fit to the mid-to-late M dwarf sample. |
The fits are made to he bius by inchicding Poisson uucertaiuties which are not shown in the plot or clarity. | The fits are made to the bins by including Poisson uncertainties which are not shown in the plot for clarity. |
It can be seeu hat the fit to the carly M stars drops wach more rapidly than the fit to tie later Ms The power aws are described by ὃν ‘Oc Sins Re 515 for the carly Ms. whereas for the later Msit is ouly found obe x. H7, higbliebhtiie the differing steepness of cach slope. | It can be seen that the fit to the early M stars drops much more rapidly than the fit to the later Ms. The power laws are described by $\delta$ $\delta$$v$ sin $i$ $\propto$ $x$$^{-3.13}$ for the early Ms, whereas for the later Ms it is only found to be $\propto$ $x$$^{-1.12}$, highlighting the differing steepness of each slope. |
The carly Ms have a uch lounger ail than the later Ms due to his faster decay of the distribution aud js is also probably au uuderestimate since we mcelude all upper linits in the data te 1lcrease he sauple size. wuch iuchided a imuber of stars from the ? sample with detection limits of LO kins 1. | The early Ms have a much longer tail than the later Ms due to this faster decay of the distribution and this is also probably an underestimate since we included all upper limits in the data to increase the sample size, which included a number of stars from the \citet{stauffer86} sample with detection limits of 10 km $^{-1}$. |
We note hat he curvature « hese slopes are also affected by he activitv lietimes shown in Table 2) since hey change with sλοςba] type. | We note that the curvature of these slopes are also affected by the activity lifetimes shown in Table \ref{tab:sat} since they change with spectral type. |
The differe1ος of 2.01 41 the exponent between he two spectral samples allows us au insight iuto he efficiency of he braking mechanisin between wtiallv aud tally convecive stars. assunudnue the arecr measured rotation in the later AI star salle Is LO ue to increased line blending TOlu Hicrease molecular xuads aud iustiuuenutal resohtion arguin | The difference of 2.01 in the exponent between the two spectral samples allows us an insight into the efficiency of the braking mechanism between partially and fully convective stars, assuming the larger measured rotation in the later M star sample is not due to increased line blending from increased molecular bands and instrumental resolution arguments. |
ien ST have shown that t uajoritv of rapidly rotating mud-AL stars are nembers of the votπιο disk population. whereas he oder populatio1 tend to rotate more slowly. | \citet{delfosse98} have shown that the majority of rapidly rotating mid-M stars are members of the young disk population, whereas the older population tend to rotate more slowly. |
Tf this is indeed the case, then the bralksi uechinism is at pav in later AL stars but f1C efficiency. has dropped across the fully convective 301ndary. | If this is indeed the case, then the braking mechanism is at play in later M stars but the efficiency has dropped across the fully convective boundary. |
As menjoned earlier. the clanee iu he fek topology between the partially axd ullv couvective botudary is probably the divi18o actorac which goveris the efficiency of he wixd raking uechanisui since ? have shown hat fi convective stars produce field streugls as stro as outiallv convective stars. | As mentioned earlier, the change in the field topology between the partially and fully convective boundary is probably the driving factor which governs the efficiency of the wind braking mechanism since \citet{reiners07a} have shown that fully convective stars produce field strengths as strong as partially convective stars. |
To beter probe raking uechanisui in this fashion. i1 acldition ealimg more data. it is necessary to aSO old iu an aeaep proxy for the sample. aud decote both vounes ale old disk populations to compare these objects. | To better probe the braking mechanism in this fashion, in addition to gaining more data, it is necessary to also fold in an age proxy for the sample, and decouple both the young and old disk populations to compare these objects. |
Along with this a better understaudi18o of the activity lifetimes and how jiese. chiauge with specral type should be considered. | Along with this a better understanding of the activity lifetimes and how these change with spectral type should be considered. |
Alore maguctic field topology studies are required for later M stars which cau help to coinia if the later Ms aso have axisvinctric poloidal fields aud | More magnetic field topology studies are required for later M stars which can help to confirm if the later Ms also have axisymetric poloidal fields and |
Willams ct al. ( | Williams et al. ( |
1991) detected. cussion at a level of 1.1«107 TE 7 at the eastern cud of TDSS. | 1991) detected emission at a level of $1.4 \times 10^{20}$ H $^{-2}$ at the eastern end of TDSS. |
In Figure 6. we show our optical VR|I nuage together with the intensity map taken from Williams ct al. ( | In Figure 6, we show our optical $VR+I$ image together with the intensity map taken from Williams et al. ( |
1991). | 1991). |
The enmüssiou appears associated with TDSS and distributed iu au ellipse whose semi-uajor aud axes are 1570 aud 1070. respectively. | The emission appears associated with TDSS and distributed in an ellipse whose semi-major and semi-minor axes are $\farcs$ 0 and $\farcs$ 0, respectively. |
If we assiue that the column density within this ellipse is constant at a value of 1.1«1079 TIE cin.2, we roughly estimate that the gas mass associated with TDSS is 2«10M... | If we assume that the column density within this ellipse is constant at a value of $1.4 \times 10^{20}$ H $^{-2}$, we roughly estimate that the gas mass associated with TDSS is $2 \times 10^{7} M_{\odot}$. |
Thus. we obtain a ratio of the eas mass to optical B baud huninosity of 3.2«102A.1. for TDSS. | Thus, we obtain a ratio of the gas mass to optical $B-$ band luminosity of $3.2 \times 10^{-2} M_{\odot}/L_{\odot}$ for TDSS. |
This value is much simaller than those of Arp 1058 (Lp=LO STOEL. My Ευ... and Arp 215N (Lp=6.1ςLOSE... Afi,=ὃν10M. AMin/Lp=LOM. /£.). | This value is much smaller than those of Arp 105S $L_{B}=4.9\times 10^{8}L_{\odot}$ , $M_{\rm H{\sc i}}
=4\times 10^{8}M_{\odot}$, $M_{\rm H{\sc i}}/L_{B}=0.8
M_{\odot}/L_{\odot}$ ) and Arp 245N $L_{B}=6.1\times 10^{8}
L_{\odot}$, $M_{\rm H{\sc i}}=6\times 10^{8}M_{\odot}$, $M_{\rm H{\sc i}}
/L_{B}=1.0M_{\odot}/L_{\odot}$ ). |
Furthermore. Villchez Ielesias-Parramo (1998) detect no Πα cussion in TDSS. | Furthermore, lchez Iglesias-P\'{a}rramo (1998) detect no $\alpha$ emission in TDSS. |
This suggests the poverty of verv massive stars iu TDSS. | This suggests the poverty of very massive stars in TDSS. |
These observations sugecst that only a small amount of gas was involved in the formation of TDSS. | These observations suggest that only a small amount of gas was involved in the formation of TDSS. |
Dettoni Fasano (1993) aud Boufauti et al. ( | Bettoni Fasano (1993) and Bonfanti et al. ( |
1999) mentioned another possibility: ie. TDSS (called “TCG TOOL” in their papers) is a pre-existing galaxy falling iuto SS that is being destroved by ealaxy-ealaxy interaction. | 1999) mentioned another possibility; i.e., TDSS (called “HCG 79b1” in their papers) is a pre-existing galaxy falling into SS that is being destroyed by galaxy-galaxy interaction. |
Since the photometric propertics of TDSS are similar to those of dEs (.e.. D-V colors aud exponeutial surface brightuess profiles). their idea may remain as a possible one. | Since the photometric properties of TDSS are similar to those of dEs (i.e., B-V colors and exponential surface brightness profiles), their idea may remain as a possible one. |
However. from the metallicity of TDSS. as determined from the SED. we sugeest that TDSS is a tidallv-iuduced object consisting primarily of stars liberated from UCC του, | However, from the metallicity of TDSS, as determined from the SED, we suggest that TDSS is a tidally-induced object consisting primarily of stars liberated from HCG 79b. |
Further investigation of the mass-to-huinosity ratio and/or a more detailed investigation of the metallicity would be useful to coufiii this hypothesis. | Further investigation of the mass-to-luminosity ratio and/or a more detailed investigation of the metallicity would be useful to confirm this hypothesis. |
It also requires uunuercal simulations to make sure whether tidally object being made by galaxy interactions without secondary star formation. | It also requires numerical simulations to make sure whether tidally object being made by galaxy interactions without secondary star formation. |
From iulti-bbaud. photometry of the prominent fidal debris feature to the northeast of Sevfert’s Sextet we have obtained the following results: We couclude that TDSS is simply a passive tidal feature ike many others iu interacting galaxy svstenis; | From multi-band photometry of the prominent tidal debris feature to the northeast of Seyfert's Sextet we have obtained the following results: We conclude that TDSS is simply a passive tidal feature like many others in interacting galaxy systems. |
We also conclude that there is uo indication that TDSS will evolve into a star-foriuug tidal dwarf galaxw similar to those xeviouslv studied. | We also conclude that there is no indication that TDSS will evolve into a star-forming tidal dwarf galaxy similar to those previously studied. |
This. however. iudicates that another ype of forming dwarf galaxy without secondary star ormation through the galaxy interaction. | This, however, indicates that another type of forming dwarf galaxy without secondary star formation through the galaxy interaction. |
We would like to thauk the staff members of he Okavama Astrophysical Observatory. the KISO observatory and the UI 2.2 ia telescope for their kind assistance during our observations. | We would like to thank the staff members of the Okayama Astrophysical Observatory, the KISO observatory and the UH 2.2 m telescope for their kind assistance during our observations. |
We thank J. W. Suleutic for his IIST observation of Sevfert'* Sextet. and an snonvnious referee for several useful conuuaeuts which helped improve the paper. | We thank J. W. Sulentic for his HST observation of Seyfert's Sextet, and an anonymous referee for several useful comments which helped improve the paper. |
We also thauk Ichi Tanaka for his kind help during our observations. Richared Wainscoat and Shinki Ovabu for their useful comments ou photometric calibration. aud Daisuke Kawata for useful discussions. | We also thank Ichi Tanaka for his kind help during our observations, Richard Wainscoat and Shinki Oyabu for their useful comments on photometric calibration, and Daisuke Kawata for useful discussions. |
A part of this work was done when YT was a visiting astronomer at the IfA. University of Tawa. | A part of this work was done when YT was a visiting astronomer at the IfA, University of Hawaii. |
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