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The residuals to this [it provide a hint that additional emission is present. | The residuals to this fit provide a hint that additional emission is present. |
The image shows an extra peak of low significance located eT" (o the southeast. | The image shows an extra peak of low significance located $\sim7''$ to the southeast. |
This additional peak may be related to the slight extension seen in the single dish 1.3 mm bolometer image of IIenning et al. ( | This additional peak may be related to the slight extension seen in the single dish 1.3 mm bolometer image of Henning et al. ( |
1998). | 1998). |
Extensions along the same position angle are also visible in scattered light. in particular in the sensitive HST/STIS images of Grady et al. ( | Extensions along the same position angle are also visible in scattered light, in particular in the sensitive HST/STIS images of Grady et al. ( |
2001). where nebulosity can be seen to extend from the northwest to the southeast. | 2001), where nebulosity can be seen to extend from the northwest to the southeast. |
No line emission is detected (o a limit of 0.17 Jv per channel (36). either al the LSR velocity of the cloud. DC296.2-7.9 of 3.6 km | (Otrupcek. Hartley Wane 2000). or at the stellar LSR velocity of about -1 km +. which derives [rom the heliocentric “astrometric” radial velocity of 9+1 km ! determined from Lipparcos (Madsen 2002). | No line emission is detected to a limit of 0.17 Jy per channel $3\sigma$ ), either at the LSR velocity of the cloud DC296.2-7.9 of 3.6 km $^{-1}$ (Otrupcek, Hartley Wang 2000), or at the stellar LSR velocity of about -1 km $^{-1}$, which derives from the heliocentric “astrometric” radial velocity of $9\pm1$ km $^{-1}$ determined from Hipparcos (Madsen 2002). |
The spatially resolved observations of the J=10 emission from TW Iva allow for a check of (he basic paradigm that trace molecules are present in an extended upper laver of the disk. with overall abundances depleted with respect to molecular hydrogen by large factors by comparison with those of interstellar dark clouds (Aikawa ILerbst 1999. Willacy | The spatially resolved observations of the $^+$ J=1–0 emission from TW Hya allow for a check of the basic paradigm that trace molecules are present in an extended upper layer of the disk, with overall abundances depleted with respect to molecular hydrogen by large factors by comparison with those of interstellar dark clouds (Aikawa Herbst 1999, Willacy |
We compared. our theoretical results with those given by Cesaroni et al. ( | We compared our theoretical results with those given by Cesaroni et al. ( |
1994). that performed high resolution observations by using LRAAL Plateau de Bure Interferomoeter. | 1994), that performed high resolution observations by using IRAM Plateau de Bure Interferometer. |
Located in the cloud NCGC1333. belonging to the Perseus complex at 220 pc in distance. LRAS2 is à low-mass παν protostellar svstem. including LRASZA and LRAS2ZB. separated. by 307. | Located in the cloud NGC1333, belonging to the Perseus complex at 220 pc in distance, IRAS2 is a low-mass binary protostellar system, including IRAS2A and IRAS2B, separated by $30''$. |
An infling envelope. a cireumstellar disk ancl multiple. outflows seem to be associated. with hem. | An infalling envelope, a circumstellar disk and multiple outflows seem to be associated with them. |
The advantage of observing such nearby. objects: is hat these sources can be spatially resolved with millimetre interferometers. | The advantage of observing such nearby objects is that these sources can be spatially resolved with millimetre interferometers. |
As à result we adopted a multipoint moclel in order to reproduce the density profile of the hot corino associated with HUXS2. as reported in. Maret ct al. ( | As a result we adopted a multipoint model in order to reproduce the density profile of the hot corino associated with IRAS2, as reported in Maret et al. ( |
2004). | 2004). |
We also used the temperature profile for hot corino given in Awad et al. ( | We also used the temperature profile for hot corino given in Awad et al. ( |
2010). | 2010). |
We evaluated the total column density for the methyl formate as cescribed in. Viti et al. ( | We evaluated the total column density for the methyl formate as described in Viti et al. ( |
2004b): where X is the methyl formate fractional abundances and el. is the visual extinction at length L (pc) | 2004b): where X is the methyl formate fractional abundances and $A_{\rm v}$ is the visual extinction at length L (pc). |
The protostar associated with the Bi-b core is believed to be between a pre-stellar anc Class 0 protostar evolutionary stage. | The protostar associated with the B1-b core is believed to be between a pre-stellar and Class 0 protostar evolutionary stage. |
“Pwo sources are identified with the Bi-b core. designed as BI-bN and BI-bS. separated by 20” in the north-south direction at about 350 pe in distance (Hirano 1999). | Two sources are identified with the B1-b core, designed as B1-bN and B1-bS, separated by $20''$ in the north-south direction at about 350 pc in distance (Hirano 1999). |
The physical parameters of two sources are very similar (Pins 18 Ix. mass of 1.6-1.8 solar masses and. luminosity about 2.6-3.1 solar luminosities). | The physical parameters of two sources are very similar $_{dust}$ 18 K, mass of 1.6-1.8 solar masses and luminosity about 2.6-3.1 solar luminosities). |
We qualitatively moclel the Bl-b core running our two-phases nioclel: for this core. during Phase LE the temperature only rises up to 30 Ix: this implies that onlv the weakly bound species partly sublimate due to thermal desorption (see Collines et al. ( | We qualitatively model the B1-b core running our two-phases model; for this core, during Phase II the temperature only rises up to 30 K; this implies that only the weakly bound species partly sublimate due to thermal desorption (see Collings et al. ( |
20042): hence any methyl formate formed. on the grains would remain in solid. phase. | 2004)); hence any methyl formate formed on the grains would remain in solid phase. |
However. non thermal desorption processes are known to be ellicient even at LO Ix (see Roberts et al. ( | However, non thermal desorption processes are known to be efficient even at 10 K (see Roberts et al. ( |
2007): Boland de Jong (1982): HLasegawa Llerbst (1093) and Drnga Johnson (2004)). | 2007); Boland de Jong (1982); Hasegawa Herbst (1993) and Bringa Johnson (2004)). |
We have therefore ran Phase | including some non thermal desorption mechanisms (as in ltoberts et al. ( | We have therefore ran Phase I including some non thermal desorption mechanisms (as in Roberts et al. ( |
2007)). | 2007)). |
We compare our results for methyl formate with those reported in Obere et al. ( | We compare our results for methyl formate with those reported in Oberg et al. ( |
2010). | 2010). |
Note that our computed. methyl formate is a lower limit as our non thermal desorption mechanisms did not include direct UV photodesorption. | Note that our computed methyl formate is a lower limit as our non thermal desorption mechanisms did not include direct UV photodesorption. |
Table 2 reports all our model results and compares the theoretical. predicted: column densities for methyl formate with those observed in the sources described: above. | Table 2 reports all our model results and compares the theoretical predicted column densities for methyl formate with those observed in the sources described above. |
A relatively good agreement within a factor of LO was found for the dark core associated with Bl. but the theoretical values are too low for the cases of warmer star forming regions. | A relatively good agreement within a factor of 10 was found for the dark core associated with B1, but the theoretical values are too low for the cases of warmer star forming regions. |
In fact around protostars the Dux of ions due to stellar Hares should. also be considered. | In fact around protostars the flux of ions due to stellar flares should also be considered. |
As discussed by Ciarozzo et al. ( | As discussed by Garozzo et al. ( |
2011). the elfects induced by cosmic rays on icy &rain mantles during the collapse phase cra are comparable to the effects induced by stellar Hare ions during the warm-up phase. | 2011), the effects induced by cosmic rays on icy grain mantles during the collapse phase era are comparable to the effects induced by stellar flare ions during the warm-up phase. |
lo summary. using the rate coellicient derived. fron laboratory experiments of ice irradiation. we investigated the viability of a cold. solid. state path to form. methyl formate during the evolution. of protostellar cores when the temperatures are low (10 Ix). | In summary, using the rate coefficient derived from laboratory experiments of ice irradiation, we investigated the viability of a cold solid state path to form methyl formate during the evolution of protostellar cores when the temperatures are low (10 K). |
We find that we cannot reproduce the abundance of methyl formate found in. hot cores and hot corinos without invoking gas- or solid-phase reactions which necessitate high temperatures (and that can therefore only oecur during the warm up phase once the star is formed). while we were able to reproduce the observed abundances in dark clouds. | We find that we cannot reproduce the abundance of methyl formate found in hot cores and hot corinos without invoking gas- or solid-phase reactions which necessitate high temperatures (and that can therefore only occur during the warm up phase once the star is formed), while we were able to reproduce the observed abundances in dark clouds. |
This work supports the idea that cosmic ion irradiation of icy grain mantles may be able to contribute to the production of the methyl formate observed in dense molecular clouds. but that other routes of formation are also recuirecd. | This work supports the idea that cosmic ion irradiation of icy grain mantles may be able to contribute to the production of the methyl formate observed in dense molecular clouds, but that other routes of formation are also required. |
The research. leacing to these results has received. funding from the European Communitys]. Seventh Framework Programme FP7/2007-2013] under grant agreement ni 238258. | The research leading to these results has received funding from the [European Community's] Seventh Framework Programme [FP7/2007-2013] under grant agreement $^{\circ}$ 238258. |
We thank an anonymous referee. for helping us improve the original manuscript. | We thank an anonymous referee for helping us improve the original manuscript. |
↴∖↴↑↸∖∐⋜∐⋅≼∐↴∖↴↨↘↽↴∖↴∐⋜↕↴∖↴↴⋝↸∖↸∖∐↕∪∏↕≼⊔∐⋜↧↴∖↴↑↸∖⋜∥∐↕⋅↖↽∐⊔⊳↥⋅↸∖⋜↧↴∖↴∐↓∶↴∙∙∙ ποιο featuresandor can.inspection ∱⋅∐∷∖↴↑≼∖↖⇁≼∖↨⋅≼∐∖↾≼∖↖⋡↾↕⊓↕↓∩↕⊳↴∖↴↻∏⋅⋮∐≓↕∐∖↽≼∖↴∖↴⊓⋅⊓⊳⋯⋅≼∖↕∐↑∖↖⇁∩ Siuuilarly. the existence of previously undetected isks in chwart elliptical (dE) ealaxies are now veine realized. | Similarly, the existence of previously undetected disks in dwarf elliptical (dE) galaxies are now being realized. |
Although it has been known for nearly tweuty vears that some dE-like ealaxics ∪∐⋜↧↖↽↸∖≼∐↴∖↴↘↽≓↕∐↘↽↸∖⋯∪↥⋅≻∐∪↕∪∶↴⋁↕↸∖↴∖↴≺≼↧≋∩∶≋⋜⋯≼↧⋜↧∶↴∙⊾↸∖∙∖↽ ↕≩↕↕↓⋮↰∙⋡⋮↰∙⋡≼∖∐↕⋂⊲∖∖⋮↕≩↕↕↓⋮↰∙⋡⋮↰∙⋡≼∖↕∙∖↽≼⊲⋮⋯↓≼∖↥⋅⊓∐↕≝↭↕≻∙↖↖↽↕⋯↾ Is new in addition. to the riiug number of disk. detectious. is. that some dE-like. galaxies. actually have (stellar) spiral structures im their isk. | Although it has been known for nearly twenty years that some dE-like galaxies do have disk-like morphologies (dS0; Sandage Binggeli 1984; Binggeli Cameron 1991), what is new — in addition to the rising number of disk detections — is that some dE-like galaxies actually have (stellar) spiral structures in their disk. |
After the initial surprise announceineut of a tightlyavound. two-arimed spiral structure in the dE otealaxy IC 3328 (Jerjen. kaluajs. Dineseli 20005). Jerjeu. Kaluajs Bingeeli (2001) and Darazza. Bineecli, Jerjen (2002) reported previously undetected spiral structure and bars (o. disks) in four more Vireo cluster ealaxies (2 dS0. 1 dE. aud 1 low-luninosity E) | After the initial surprise announcement of a tightly-wound, two-armed spiral structure in the dE galaxy IC 3328 (Jerjen, Kalnajs, Binggeli 2000b), Jerjen, Kalnajs Binggeli (2001) and Barazza, Binggeli, Jerjen (2002) reported previously undetected spiral structure and bars (i.e., disks) in four more Virgo cluster galaxies (2 dS0, 1 dE, and 1 low-luminosity E). |
De Rijeke et (2003) have additionally presented photometric and kincmatic evidence for disks. aud in one case spiral armis. in two edec-on Foruax cluster 90 ealaxies. | De Rijcke et (2003) have additionally presented photometric and kinematic evidence for disks, and in one case spiral arms, in two edge-on Fornax cluster dS0 galaxies. |
The data to date suggests 20% "iof Muοἱof -niehtlates eel⋅cwarfpegalaxies dear ↖galas | The data to date suggests that up to of bright early-type dwarf galaxies in clusters may have disks. |
tht ∙∖↽∫⇀∪∐∶↴∙⊾∪↓∩∩↓∶⊱⋯↥⋅∑⋜↧∙∖↽↕≧↸∖∐≼∐∖↥⋅↓∩∩⋅↱⊐∶≼∶↥⋅⋜↧∐⋜⋯⊔∖↑ al. carly-type galaxies residing iu the deusest 1998: euvironnient studied so far. namely the Coma cluster. | We report here on the first ever detection of spiral-like structure in two dwarf, early-type galaxies residing in the densest cluster environment studied so far, namely the Coma cluster. |
The presence. or at least detection. of (stellar) spiral patterus in dawarf galaxies is a particularly rare phenomenon. | The presence, or at least detection, of (stellar) spiral patterns in dwarf galaxies is a particularly rare phenomenon. |
Although dwarf versious of Sin and br galaxies have been known for a loug time | Although dwarf versions of Sm and Irr galaxies have been known for a long time |
fora,©»Lf2. | for $\alpha_n>-1/2$. |
Therefore. if we assume that a,21/2. it follows that Now. if (Diya,OW))w-0=0 for all j(Q.a,). then equation (34)) leads to ConsequentLy. and the DE is This equation. for a,=0. corresponds to the method developed in Walnajs(1976)... working in an adequate rotating framo. | Therefore, if we assume that $\alpha_n > -1/2$, it follows that Now, if $(D_\Psi^j\sigma_n(\Psi))_{\Psi = 0} = 0$ for all $j \in (0,\alpha_n)$, then equation \ref{newsigma4}) ) leads to Consequently, and the DF is This equation, for $\alpha_n=0$, corresponds to the method developed in \cite{kal}, working in an adequate rotating frame. |
As a particular case. suppose that Then. by taking the fractional derivative we obtain and the DE is This relation can be interpreted as the analogous case of the Fricke expansion. when we are dealing with Lat svstenis. | As a particular case, suppose that Then, by taking the fractional derivative we obtain and the DF is This relation can be interpreted as the analogous case of the Fricke expansion, when we are dealing with flat systems. |
Lt can be verified performing the pseudo-volumoe density (3)) of RAW and taking the Erike component corresponding to the tridimensional case. | It can be verified performing the pseudo-volume density \ref{seudorho}) ) of $R^{2\alpha_n}\Psi^{\beta_k}$ and taking the Frike component corresponding to the tridimensional case. |
We can generalize the result (37)) i£ we express the DE. in terms of Q—5L2/(2402). | We can generalize the result \ref{dftotal}) ) if we express the DF in terms of $Q
= \varepsilon-L_z^2/(2R_a^2)$. |
In this way. if the density has the form the corresponding DE is [or BR,70 and à,71/2. | In this way, if the density has the form the corresponding DF is for $R_a > 0$ and $\alpha_n > -1/2$. |
Furthermore. if we consider models with gravitational potential having no upper bound. we can deduce that for a density and assuming that lim,σι)=0 for all αν). then | Furthermore, if we consider models with gravitational potential having no upper bound, we can deduce that for a density and assuming that $\lim_{_{\Phi \rightarrow \infty}} D_\Phi^j \sigma_n(\Phi) =0$ for all $j\in(0,\alpha_n)$ , then for $R_a > 0$, $\alpha_n > -1/2$ and $Q = E + L_z^2/(2R_a^2)$. |
In this section we will use the formulae introduced above to the Binnes’s logarithmic model and the Alestel disc. and we will see that their corresponding DEs match exactly with those that were found through the application of other methods. | In this section we will use the formulae introduced above to the Binney's logarithmic model and the Mestel disc, and we will see that their corresponding DFs match exactly with those that were found through the application of other methods. |
Binney’s logarithmic model has a gravitational potential of the form. whereas its mass density is that can be written as as D? cSate’for any a €πιwe obtain the same DE founded in Jiang&Ossipkov(2007).. using the Abel's integral equation. and in Evans(1993b) using Lynden-Bell’s method. | Binney's logarithmic model has a gravitational potential of the form whereas its mass density is that can be written as So, as $D_x^\alpha e^{ax}=a^\alpha e^{ax}$ for any $\alpha\in\mathbb{R}$, we obtain where the same DF founded in \cite{jiang}, using the Abel's integral equation, and in \cite{eva93b} using Lynden-Bell's method. |
Another case of interest is the Moestel disc. characterized by a gravitational potentialof the form and a surface mass density given by where X,.=erDorey(οπέςη). | Another case of interest is the Mestel disc, characterized by a gravitational potentialof the form and a surface mass density given by where $\Sigma_0=v_c^2/(2 \pi G R_0)$. |
Now.. we can write. (50)) as [or any m€Iz. | Now, we can write \ref{denmestel}) ) as for any $m\in\mathbb{R}$. |
So. equation (44)) for Ry=o Leads to where £ and e are the constants given by "his solution was obtained as well by Evans (1993b).. and was proposed earlier by Toomre (1971).. | So, equation \ref{nouppf}) ) for $R_a\rightarrow\infty$ leads to where $F$ and $\sigma$ are the constants given by This solution was obtained as well by \cite{eva93b}, , and was proposed earlier by \cite{T3}. . |
The main purpose of this work las beeu to find a 1nore effective compression iiethod for floatine-point astronomical dmaees than is provided by the ↕∪↴∖↴↴∖↴↕↸∖↴∖∷∖↴⋯↸∖↑∐∪≼↧↴∖↴↑∐⋜↧↑⋜⋯∖↸⊳∪⋯⋯∪↕∐⋅↖↽∏↴∖↴↸∖≼↧≺↴∖↴⋯⊳∐⋜↧↴∖↴ GZIP). | The main purpose of this work has been to find a more effective compression method for floating-point astronomical images than is provided by the lossless methods that are commonly used (such as GZIP). |
These Soatiug-poiut+ images+ typically+ do not compress well with lossless algorithms because a large fraction of the bits in cach pixel value representation contain no significant information and are effectively filled with uuconipressible noise. | These floating-point images typically do not compress well with lossless algorithms because a large fraction of the bits in each pixel value representation contain no significant information and are effectively filled with uncompressible noise. |
We have adopted the method of elininatiug some of this noise by quantizing the pixel values into a set of discrete. linearly spaced intensity levels. which are represented by scaled integer values. | We have adopted the method of eliminating some of this noise by quantizing the pixel values into a set of discrete, linearly spaced intensity levels, which are represented by scaled integer values. |
The scaled integers cau then be efficieutlv conrpressed using the very fast Rice algoritlun. | The scaled integers can then be efficiently compressed using the very fast Rice algorithm. |
Iu order to make these compression techniques nore widely available. we have produced a pair of utility programs calledfpack aud that cau ve used to compress auv FITS format tage. | In order to make these compression techniques more widely available, we have produced a pair of utility programs called and that can be used to compress any FITS format image. |
For convenience. we define a quantization yavaancter. q. which is equal to the measted RMS io]se In background regions of the mage divided w the spacing between the quantized intensity evels. | For convenience, we define a quantization parameter, q, which is equal to the measured RMS noise in background regions of the image divided by the spacing between the quantized intensity levels. |
Given a particular q value when quantizing and compressing an dniage. one can calculate roni fundamental principles the expected mage conrpression ratio. as eiven by equation 6.. | Given a particular q value when quantizing and compressing an image, one can calculate from fundamental principles the expected image compression ratio, as given by equation \ref{eq:ratio3}. |
Coarser quantization (ic.. smaller q values] gives ercater nuage compression. but at the same time it increases the RAIS pisel-to-pixcl noise bv an amount eiven by equation LO which also teuds o degrade the precision of the magnitude aud »ositiou mnieasurements of the objects in the nage. | Coarser quantization (i.e., smaller q values) gives greater image compression, but at the same time it increases the RMS pixel-to-pixel noise by an amount given by equation \ref{eq:fractionalnoise} which also tends to degrade the precision of the magnitude and position measurements of the objects in the image. |
Qur series of experiments on sinulated and ou real astronomical CCD images demonstrate that he noise equation 10 also gives a good estimate of the iucrease in measurement uncertainties or objects near the detection threshold in a quantized iuage (which are noie limited). | Our series of experiments on simulated and on real astronomical CCD images demonstrate that the noise equation \ref{eq:fractionalnoise} also gives a good estimate of the increase in measurement uncertainties for objects near the detection threshold in a quantized image (which are noise limited). |
For uanyv practical applications. q values between Lo aud 1 provide a good combination of lieh conrpression and low imereased noise: q = 1 gives a colupression ratio of about 10 while increasing the noise and the measurement uncertainties by about aud q = L gives a compression ratio of 6 with a neeligible increase. iu. the noise. aud iineasurement iuicertaiuties. | For many practical applications, q values between 1 and 4 provide a good combination of high compression and low increased noise: q = 1 gives a compression ratio of about 10 while increasing the noise and the measurement uncertainties by about, and q = 4 gives a compression ratio of 6 with a negligible increase in the noise and measurement uncertainties. |
. Iu "quick-look types of applications. where high | In “quick-look” types of applications, where high |
sale loops are more potential aud less stretched in Case C as seen in the bottom panel of Figure ὃν, | same loops are more potential and less stretched in Case C as seen in the bottom panel of Figure \ref{fig:f3}. |
The interplay between the radial speed and the coronal density structure determines the stellar niass loss rate. as well as the stellar angular moineutuu loss rate to the stellar wind. | The interplay between the radial speed and the coronal density structure determines the stellar mass loss rate, as well as the stellar angular momentum loss rate to the stellar wind. |
We follow the method by Cohenetal.(2009).. and calculate these loss rates from the MIID solution. | We follow the method by \cite{Cohen09b}, and calculate these loss rates from the MHD solution. |
This method expands the idealized approach bv Weber&Davis(1967) aud uses the fact that the ΑΠΙΟ solution provides a realistic. non-idealized Alfvén surface. at which the Alfvéunic Mach umber. M4=ufey= dl. where οι=BfTrp is the Alfvéun speed. | This method expands the idealized approach by \cite{weberdavis67} and uses the fact that the MHD solution provides a realistic, non-idealized Alfvénn surface, at which the Alfvénnic Mach number, $M_A=u/v_A=1$ , where $v_A=B/\sqrt{4\pi\rho}$ is the Alfvénn speed. |
Ouce the Alfvéóun surface has been determined. the loss rates can be calculated as: where ra is the local radius of the Alfvémn surface. da, is a surface element. aud the iutegration is done over the realistic Alfvénn surface. | Once the Alfvénn surface has been determined, the loss rates can be calculated as: where $r_A$ is the local radius of the Alfvénn surface, $\mathbf{da}_A$ is a surface element, and the integration is done over the realistic Alfvénn surface. |
It is worth to mentioning that the realistic Alfvén surface. at which the maguctic breaking of the stellar wind takes place. is the actual source surface. and it does not have a spherical shape as is assuned to have in the couunon use of the potential field approximation. | It is worth to mentioning that the realistic Alfvénn surface, at which the magnetic breaking of the stellar wind takes place, is the actual source surface, and it does not have a spherical shape as is assumed to have in the common use of the potential field approximation. |
The mass aud angular momentum loss rates for the different test cases are shown in the upper part of Table 2.. | The mass and angular momentum loss rates for the different test cases are shown in the upper part of Table \ref{table:t2}. |
For comparison aud verification of these results. we have computed simular wind models for the solar case. | For comparison and verification of these results, we have computed similar wind models for the solar case. |
The same σα] wethod described above was cluploved for a solar αποmaguctoeram obtained durius the last solar maxim (Carrineton Rotation 1958) by the SoTO | The same numerical method described above was employed for a solar magnetogram obtained during the last solar maximum (Carrington Rotation 1958) by the SoHO. |
This solar simulation resulted ii a iss loss vate of MIDI’...Af,=2.10HAF.Vr| and an angular moment loss rate of Fx1009g0n?s7 with the use of base density of 2«105 3 (same as Case A). | This solar simulation resulted in a mass loss rate of $\dot{M}_\odot\approx 2\cdot10^{-14}\;M_\odot\;Yr^{-1}$ and an angular momentum loss rate of $\dot{J}\approx 10^{30}\;g\;cm^2\;s^{-2}$ with the use of base density of $2\times 10^{8}$ $^{-2}$ (same as Case A). |
These are sinular to canonical solar values. as expected. | These are similar to canonical solar values, as expected. |
Iustead. the loss rates πο AB Dor are significantly lugher than solar. | Instead, the loss rates from AB Dor are significantly higher than solar. |
The augular momeutua loss rate cau be 2 orders of magnitude higher. simply due to the much more rapid rotation of AB Dor. | The angular momentum loss rate can be 2 orders of magnitude higher, simply due to the much more rapid rotation of AB Dor. |
For Case A. the mass loss rate for AB Dor is about a factor of 10 larger thau the equivalent solar case. | For Case A, the mass loss rate for AB Dor is about a factor of 10 larger than the equivalent solar case. |
This is perhaps slightly surprising siuce all parameters other than the rotation rate aud. to sole extent. the surface field map are fairly similar to those of the active Sun chosen for the comparison. | This is perhaps slightly surprising since all parameters other than the rotation rate and, to some extent, the surface field map are fairly similar to those of the active Sun chosen for the comparison. |
The most conspicuous difference is the factor of 50 in rotation rate and it is worth examining the influence of rotation alone in more detail. | The most conspicuous difference is the factor of 50 in rotation rate and it is worth examining the influence of rotation alone in more detail. |
We repeated the AB Dor computational rims for Cases A-C for a rotation period of 254 instead of 0.54. with all other aspects of the simulations remaining the same. | We repeated the AB Dor computational runs for Cases A-C for a rotation period of $25\;d$ instead of $0.5\;d$, with all other aspects of the simulations remaining the same. |
The solutions for these uus are shown in Figure 9.. | The solutions for these runs are shown in Figure \ref{fig:f4}. |
The azimuthal taneling of the coronal field that characterizes the 0.5d period results in Fieure 8 is. wusurprisinely. conipletelv abseut im this set of solutious and all field lines are esseutiallv racial. | The azimuthal tangling of the coronal field that characterizes the $0.5\;d$ period results in Figure \ref{fig:f3} is, unsurprisingly, completely absent in this set of solutions and all field lines are essentially radial. |
In addition. more field. lines are open iu these solutions compared to the case with rapid rotation. | In addition, more field lines are open in these solutions compared to the case with rapid rotation. |
Mass and angular monieutum loss rates are listed at the bottom of Table 2: mass loss rates are typically a factor of teu lower than for the 0.54 period results. aud. for the Case A base density. are more siuilar to the solar value. | Mass and angular momentum loss rates are listed at the bottom of Table \ref{table:t2}: mass loss rates are typically a factor of ten lower than for the $0.5\;d$ period results, and, for the Case A base density, are more similar to the solar value. |
The MID. simulations of AB Dor preseuted here reveals a coronal structure that is manifestly differcu roni the woelbstudied solar corona. | The MHD simulations of AB Dor presented here reveals a coronal structure that is manifestly different from the well-studied solar corona. |
These differences are due to the different magnetic structure. which is uostlv composed of high-latitude. large-scale regious of strong imnagnetie field. aud the rapid stellar rotation which induces aziuauthal wrapping and taneling of nagnetic feld. | These differences are due to the different magnetic structure, which is mostly composed of high-latitude, large-scale regions of strong magnetic field, and the rapid stellar rotation which induces azimuthal wrapping and tangling of magnetic field. |
This taugliug cannot be obtained from he static. non-MIID. potential field extrapolation. which is eenerally useful ouly for studving the small close oops near the surface where elobal effects are less imuportaut. | This tangling cannot be obtained from the static, non-MHD, potential field extrapolation, which is generally useful only for studying the small closed loops near the surface where global effects are less important. |
We note in passing that for stars with larec-scale regions of strong magnetic feld. closed loops are xobablv siguificautlv larecr than in the solar case aud he choice for the location of source surface location should be at ereater radial distance thin the como use of Ry=διBSR, | We note in passing that for stars with large-scale regions of strong magnetic field, closed loops are probably significantly larger than in the solar case and the choice for the location of source surface location should be at greater radial distance than the common use of $R_{ss}=2.5-3.5R_\star$. |
The simulation results show that the mass loss and angular ΠΟΙΟΙΤΗ loss rates increase with increasing coronal base density. | The simulation results show that the mass loss and angular momentum loss rates increase with increasing coronal base density. |
The explanation for the former is trivial: introducing a greater mass source at the base will necessarily increase the mass flux through a closed surface around this source. | The explanation for the former is trivial: introducing a greater mass source at the base will necessarily increase the mass flux through a closed surface around this source. |
The latter effect is ore subtle and is due to the fact that 7Xpu. | The latter effect is more subtle and is due to the fact that $\dot{J}\propto\rho u$. |
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