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However. it is not strictly obeyed. particularly if there is Ni present in the outermost regions of the SN ejecta. as discussed by ?.. | However, it is not strictly obeyed, particularly if there is Ni present in the outermost regions of the SN ejecta, as discussed by \citet{pinto00a}. |
Since Arnett's Rule relates only global quantities. or any class of SN model in which a viewing angle dependence of he flux appears (including. but not limited to. the off-centre models discussed here). an attempt to measure the nickel-mass via Arnett's Rule must introduce a systematic error which depends upon the direction of observation. | Since Arnett's Rule relates only global quantities, for any class of SN model in which a viewing angle dependence of the flux appears (including, but not limited to, the off-centre models discussed here), an attempt to measure the nickel-mass via Arnett's Rule must introduce a systematic error which depends upon the direction of observation. |
In this section. we will use the toy models to investigate how this systematic would behave if real SNe tor at least a subset of them) were to harbour lop-sided distributions of “ONL. | In this section, we will use the toy models to investigate how this systematic would behave if real SNe (or at least a subset of them) were to harbour lop-sided distributions of $^{56}$ Ni. |
First we consider the case where a Ni mass is deduced by the application of Arnett’s Rule (with a chosen value. or range of values. for à) to direct measurements of the observed peak magnitude and time of maximum light. | First we consider the case where a Ni mass is deduced by the application of Arnett's Rule (with a chosen value, or range of values, for $\alpha$ ) to direct measurements of the observed peak magnitude and time of maximum light. |
To elucidate this case. we over-plot in Figure + the curves obtained by combining equations 3 and 4 for three different values of à (0.8. 1.0. 1.2). | To elucidate this case, we over-plot in Figure \ref{fig:toy-arnett} the curves obtained by combining equations 3 and 4 for three different values of $\alpha$ (0.8, 1.0, 1.2). |
It comes as no surprise that Arnett’s Rule with a fixed value of à does not describe the co-dependence of AZ), and /,, obtained from the models — the Rule is concerned with relating global properties of different SNe. not the detailed angular dependence within single objects. | It comes as no surprise that Arnett's Rule with a fixed value of $\alpha$ does not describe the co-dependence of $M_{\rm p}$ and $t_{\rm p}$ obtained from the models – the Rule is concerned with relating global properties of different SNe, not the detailed angular dependence within single objects. |
However. the correlation of AJ, and ἐν deriving from Arnett's Rule in the same sense as that obtained from the models. | However, the correlation of $M_{\rm p}$ and $t_{\rm p}$ deriving from Arnett's Rule in the same sense as that obtained from the models. |
This helps to suppress the systematic error one would introduce with an unknown viewing angle: for the more moderate models (A and B). the systematic error introduced would be only around 0.15 to 0.2 mag. | This helps to suppress the systematic error one would introduce with an unknown viewing angle; for the more moderate models (A and B), the systematic error introduced would be only around 0.15 to 0.2 mag. |
In exceptional cases. however. a much larger systematic error of up to about 0.5 mag is possible for the adopted Ni mass. | In exceptional cases, however, a much larger systematic error of up to about 0.5 mag is possible for the adopted Ni mass. |
Secondly. we consider the case where reliable measurement of /, has not been possible so that an estimate of A; must be obtained from AZ, alone. | Secondly, we consider the case where reliable measurement of $t_{\rm p}$ has not been possible so that an estimate of $M_{\mbox{\scriptsize Ni}}$ must be obtained from $M_{\rm p}$ alone. |
Adopting the relationship discussed by ? (their equation 7 which is derived from Arnett’s Rule and an empirically motivated assumption for the value of /,. leads to an expected value of A,=—18.56c0.16 mag for 0.4 M.. | Adopting the relationship discussed by \cite{stritzinger05} (their equation 7 which is derived from Arnett's Rule and an empirically motivated assumption for the value of $t_{\rm p}$ ), leads to an expected value of $M_{\rm p} = -18.56 \pm 0.16$ mag for $M_{\mbox{\scriptsize Ni}} = 0.4$ $_{\odot}$. |
This range of A, is indicated in Figure 4. by the grey shaded region: thus all the points which fall within this band are consistent with the ? relationship. | This range of $M_{\rm p}$ is indicated in Figure \ref{fig:toy-arnett} by the grey shaded region; thus all the points which fall within this band are consistent with the \cite{stritzinger05} relationship. |
For all the models considered. a significant fraction of the possible viewing directions lead to light curves which lie outside the range. | For all the models considered, a significant fraction of the possible viewing directions lead to light curves which lie outside the range. |
Thus. if the range of light curve properties produced by the models were representative of those from real SNe. using that relationship between Ad, and Ads; in the analvsis of a sample of observed light curves could lead one to infer a wider spread of nickel masses than required. | Thus, if the range of light curve properties produced by the models were representative of those from real SNe, using that relationship between $M_{\rm p}$ and $M_{\mbox{\scriptsize Ni}}$ in the analysis of a sample of observed light curves could lead one to infer a wider spread of nickel masses than required. |
The tov models have demonstrated that it is possible to construct very simplistie scenarios in which a lop-sided distribution of "Ni can introduce significant angular dependence in the emergent radiation. | The toy models have demonstrated that it is possible to construct very simplistic scenarios in which a lop-sided distribution of $^{56}$ Ni can introduce significant angular dependence in the emergent radiation. |
The scale of the variation is comparable to that introduced by other types of departure from spherical symmetry (2:: 2)) and can be comparable to. or larger than. typical observational uncertainties. | The scale of the variation is comparable to that introduced by other types of departure from spherical symmetry \citealt{hoeflich91}; \citealt{kasen04}) ) and can be comparable to, or larger than, typical observational uncertainties. |
However. the toy models are very simplistic and are derived with no reference to realistic explosion physies. | However, the toy models are very simplistic and are derived with no reference to realistic explosion physics. |
Therefore. they merely illustrate possible effects and one must appeal to 3D explosion models to judge whether such effects are likely to have a part to play in reality. | Therefore, they merely illustrate possible effects and one must appeal to 3D explosion models to judge whether such effects are likely to have a part to play in reality. |
The remainder of this paper is concerned with the analysis of results from one such model. | The remainder of this paper is concerned with the analysis of results from one such model. |
The 3D explosion simulation 3T2d200 deseribed by ?. has been used as the basis for a model to explore the effects of an off-centre "Ni distribution in a realistic case. | The 3D explosion simulation 3T2d200 described by \citet{roepke07} has been used as the basis for a model to explore the effects of an off-centre $^{56}$ Ni distribution in a realistic case. |
This simulation followed the flame evolution when ignited in a lop-sided teardrop-like shape (upper left panel of Figure 51). | This simulation followed the flame evolution when ignited in a lop-sided teardrop-like shape (upper left panel of Figure \ref{fig:td}) ). |
Such an ignition configuration is motivated by recent studies of the pre-ignition convective burning phase (2)... which suggest that the flow pattern is dominated by a dipole at this stage. | Such an ignition configuration is motivated by recent studies of the pre-ignition convective burning phase \citep{kuhlen06}, which suggest that the flow pattern is dominated by a dipole at this stage. |
A consequence could be a lop-sided ignition as adopted here. where the majority of ignition kernels are locatec on one side of the star but with a slight over-shooting across the centre. | A consequence could be a lop-sided ignition as adopted here, where the majority of ignition kernels are located on one side of the star but with a slight over-shooting across the centre. |
In this case. the flame propagates in both directions. subjec o buovaney instabilities and accelerated by turbulence. | In this case, the flame propagates in both directions, subject to buoyancy instabilities and accelerated by turbulence. |
One side of the flame structure. however. remains dominant (see the top lef xinel in Figure 55). | One side of the flame structure, however, remains dominant (see the top left panel in Figure \ref{fig:td}) ). |
Although the ash bubble on this side of the star starts to sweep around the core (similar to what has been describec by 2)) it does not collide on the far side (see the panels in the second row of Figure 51) because the energy generated by nuclear burning is sufficient to gravitationally unbind the star and expand i before a collision occurs. | Although the ash bubble on this side of the star starts to sweep around the core (similar to what has been described by \citealt{plewa04}) ), it does not collide on the far side (see the panels in the second row of Figure \ref{fig:td}) ) because the energy generated by nuclear burning is sufficient to gravitationally unbind the star and expand it before a collision occurs. |
The lower left panel of Figure 5. shows he end-stage of this evolution: subsequently the expansion is approaching homology. | The lower left panel of Figure \ref{fig:td} shows the end-stage of this evolution; subsequently the expansion is approaching homology. |
The composition of the ejecta is illustratec in the lower right panel of Figure 5.. showing a pronounced asymmetric bubble of iron-group elements. | The composition of the ejecta is illustrated in the lower right panel of Figure \ref{fig:td}, showing a pronounced asymmetric bubble of iron-group elements. |
The simulation was carried out on a 512? cells moving Cartesian computational eric (2). | The simulation was carried out on a $^3$ cells moving Cartesian computational grid \citep{roepke05b}. |
The model used here for radiation transport adopts a uniform Cartesian grid with 1287 cells. | The model used here for radiation transport adopts a uniform Cartesian grid with $^3$ cells. |
The total mass density and the mass fraction of Fe-group elements in each cell was obtained directly from the hydrodynamic simulation by combining sets of 4° cells from the original 512 grid. | The total mass density and the mass fraction of Fe-group elements in each cell was obtained directly from the hydrodynamic simulation by combining sets of $^3$ cells from the original $^3$ grid. |
The cell properties were extracted from the latest time to which the hydrodynamics calculation was followed (~ 10s). | The cell properties were extracted from the latest time to which the hydrodynamics calculation was followed $\sim$ 10s). |
For times beyond this point. including the entirety of the time for which the radiative transfer is followed. the ejecta are assumed to be in homologous expansion. | For times beyond this point, including the entirety of the time for which the radiative transfer is followed, the ejecta are assumed to be in homologous expansion. |
Since the hydrodynamical results did not provide detailed information on the composition of the ash material. it has been assumed that initially allthe Fe-group mass was composed of "Ni: this gives a total of 0.448 M. of “Ni. | Since the hydrodynamical results did not provide detailed information on the composition of the ash material, it has been assumed that initially allthe Fe-group mass was composed of $^{56}$ Ni; this gives a total of 0.448 $_{\odot}$ of $^{56}$ Ni. |
Since this investigation is concerned with the viewing-direction dependence of the light curve. this assumption will not signiticantly affect our conclusions —alower mass fraction of "Ni in the ash would merely reduce the | Since this investigation is concerned with the viewing-direction dependence of the light curve, this assumption will not significantly affect our conclusions –alower mass fraction of $^{56}$ Ni in the ash would merely reduce the |
any additional X-ray source with B<19 to be identified with respect to 2MASS: B=19 corresponds to a star in NGC 752 of less than 0.45Mg (according to the isochrones of Girardietal.. 2002)). | any additional X-ray source with $B \la 19$ to be identified with respect to 2MASS; $B =19$ corresponds to a star in NGC 752 of less than $0.45~M_{\sun}$ (according to the isochrones of \citealp{gbb+02}) ). |
We expect most of the objects with B>19.5 to be extragalactic and. in addition. if a few of them were members of NGC 752. their mass would be well below the target of this study. which are solar-mass stars with 0.8XMS1.2Mo. | We expect most of the objects with $B > 19.5$ to be extragalactic and, in addition, if a few of them were members of NGC 752, their mass would be well below the target of this study, which are solar-mass stars with $0.8 \la M \la 1.2 M_{\sun}$. |
Accordingly. we did not further consider the catalogue to identify potential members of NGC 752 among the ssourees. | Accordingly, we did not further consider the catalogue to identify potential members of NGC 752 among the sources. |
It is interesting. however. to look at the scatter plot of the X-ray counts versus B magnitudes for the X-ray sources with a counterpart shown in reffig:distB.. | It is interesting, however, to look at the scatter plot of the X-ray counts versus $B$ magnitudes for the X-ray sources with a counterpart shown in \\ref{fig:distB}. |
A lack of sources with 14«B18 ts clearly noticeable and does not appear to be due to the sensitivity of the ddata. (as weaker sources are detected for other values of B). | A lack of sources with $14 < B < 18$ is clearly noticeable and does not appear to be due to the sensitivity of the data, (as weaker sources are detected for other values of $B$ ). |
This indicates that the lack of X-ray sources with à counterpart with 14<B18 is intrinsic to the field. consistent with the observation that NGC 752 appears to be deficient in low mass stars. | This indicates that the lack of X-ray sources with a counterpart with $14 < B < 18$ is intrinsic to the field, consistent with the observation that NGC 752 appears to be deficient in low mass stars. |
We estimated the number of expected background sources in the observation using the logN—S relationship for X-ray sources in Brandtetal.(2001). | We estimated the number of expected background sources in the observation using the $\log N-\log S$ relationship for X-ray sources in \cite{bah+01}. |
. As in 66. the faintest source has a typical count rate of 0.035::assuming for extragalactic sources a power-law spectrum with spectral index LU=L4 and a total interstellar absorption column of 5x10 em™ (Dickey&Lockman.1990: Kalberlaet 20051). this count rate to a flux of 2.2 x107! iin the energy band 0.5—2 keV. In this energy band and at this flux level. the expected number density of background sources determined on the basis of the logN—$ relationship given in Brandtetal.(2001) is 2100-2610 per square degree. | As in 6, the faintest source has a typical count rate of 0.035;assuming for extragalactic sources a power-law spectrum with spectral index $\Gamma=1.4$ and a total interstellar absorption column of $N({\rm H})
= 5 \times 10^{20}$ $^{-2}$ \citealp{dl90}; \citealp{kbh+05}) ), this count rate to a flux of 2.2 $\times 10^{-16}$ in the energy band $0.5-2$ keV. In this energy band and at this flux level, the expected number density of background sources determined on the basis of the $\log N-\log S$ relationship given in \cite{bah+01} is 2100-2610 per square degree. |
The area imaged by lis 0.08 square degree. so that the expected number of serendipitous extra-galactic sources 1s between 170 and 210. consistent with the hypothesis that most of non-identified sources are extragalactic. | The area imaged by is 0.08 square degree, so that the expected number of serendipitous extra-galactic sources is between 170 and 210, consistent with the hypothesis that most of non-identified sources are extragalactic. |
For the brightest sources. for which a spectral and timing analysis was carried out. source and background regions were defined in ps9.. and light curves and spectra were extracted | For the brightest sources, for which a spectral and timing analysis was carried out, source and background regions were defined in , and light curves and spectra were extracted |
under realistic assumptions about the total mass. the orbits would be reasonable lor a GC, with separations ranging [rom a few to 1020 AU. | under realistic assumptions about the total mass, the orbits would be reasonable for a GC, with separations ranging from a few to 10–20 AU. |
These BSSs do not show anv evidence of photometric variability. (Ixaluzny. et al. | These BSSs do not show any evidence of photometric variability (Kaluzny et al. |
1997). and variations of Via wilh full amplitudes exceeding 3kms+ on a time interval of 72 hours are excluded by our observations for BSSs it442424 and 4664677 (no such information is available for the [ast rotators). | 1997), and variations of $V_{\rm rad}$ with full amplitudes exceeding $3\kms$ on a time interval of 72 hours are excluded by our observations for BSSs 42424 and 64677 (no such information is available for the fast rotators). |
However. these results do not disprove that they. are in binary svstems. | However, these results do not disprove that they are in binary systems. |
For instance. (hey. are still consistent with what expected for ~90% of binaries characterized by an eccentricitv-period distribution similar to that recently observed by Mathieu Geller (2009) and populating the tail of the velocity distzibution in M4 (Mathieu Geller. private communication). | For instance, they are still consistent with what expected for $\sim 90$ of binaries characterized by an eccentricity-period distribution similar to that recently observed by Mathieu Geller (2009) and populating the tail of the velocity distribution in M4 (Mathieu Geller, private communication). |
Indeed. further observations are urged to search [or clear-cut signatures of binaritv. | Indeed, further observations are urged to search for clear-cut signatures of binarity. |
The fact that none of the 14 DSSs for which we measured C and/or O abundances shows signatures of depletion is quite intriguing. | The fact that none of the 14 BSSs for which we measured C and/or O abundances shows signatures of depletion is quite intriguing. |
Out of the 42 BSSs investigated in 47Tuc. FOG found that 6 (1454)) are C-depleted. with 3 of them also displaving O-depletion. | Out of the 42 BSSs investigated in 47Tuc, F06 found that 6 ) are C-depleted, with 3 of them also displaying O-depletion. |
Accordinglv. in M4 we could have expected 1-2 BsSs with depleted carbon abundance. ancl 0-1 BSS with both Ci and ο depletion. | Accordingly, in M4 we could have expected 1-2 BSSs with depleted carbon abundance, and 0-1 BSS with both C and O depletion. |
Hence. the resulting non-detection may. just be an effect of low statistics and is still consistent with the expectations. | Hence, the resulting non-detection may just be an effect of low statistics and is still consistent with the expectations. |
Alternatively the lack of chemical anomalies in M4 DSSs might point to a different Formation process: while at least 6 DS5s (the CO-depleted ones) in 47Tuc display surface abundances consistent with the MT formation channel. all the investigated) BSSs in M4 may derive from stellar collisions. for. which no chemical anonmalees are expected. | Alternatively the lack of chemical anomalies in M4 BSSs might point to a different formation process: while at least 6 BSSs (the CO-depleted ones) in 47Tuc display surface abundances consistent with the MT formation channel, all the (investigated) BSSs in M4 may derive from stellar collisions, for which no chemical anomalies are expected. |
Finally. it is also possible that the CO-depletion is a transient phenomenon (FOG) and (at least part of) the D55s in AL are indeed. MT-D55Ss which already evolved back to normal chemical abundances. | Finally, it is also possible that the CO-depletion is a transient phenomenon (F06) and (at least part of) the BSSs in M4 are indeed MT-BSSs which already evolved back to normal chemical abundances. |
The most intriguing result of this study is the discovery that a large fraction. (40%)) of the investigated. D55s in M4 are fast rotators. with rotational velocities ranging from ~50kms'. up to more than 150kms'. | The most intriguing result of this study is the discovery that a large fraction ) of the investigated BSSs in M4 are fast rotators, with rotational velocities ranging from $\sim 50\kms$, up to more than $150\kms$. |
We emphasize that this is the largest population of [ast rotating BSSs ever found in a cluster. | We emphasize that this is the largest population of fast rotating BSSs ever found in a cluster. |
Approximately 30% of the DSSs spinning faster (at 20—50kms. 1) than MS stars of the same colour has been recently found in the old open cluster NGCISS (Mathieu Geller 2009). while ος in vounger open clusters are found to rotate slower than expected for their spectral type (e.g.. Shetrone sandquist 2000; Schónnberner et al. | Approximately $30\%$ of the BSSs spinning faster (at $20-50\kms$ ) than MS stars of the same colour has been recently found in the old open cluster NGC188 (Mathieu Geller 2009), while BSSs in younger open clusters are found to rotate slower than expected for their spectral type (e.g., Shetrone Sandquist 2000; Schönnberner et al. |
2001). | 2001). |
For GCs only scarce ancl sparse data have been collected to date. | For GCs only scarce and sparse data have been collected to date. |
The most studied case is that of ArTuc. where 3 (7%)) BSSs out of the 45 measured objects (Shara et al. | The most studied case is that of 47Tuc, where 3 ) BSSs out of the 45 measured objects (Shara et al. |
1997: De Marco et al. | 1997; De Marco et al. |
2005: F06) have rotational velocities larger (han 50kms|. up to ~155kms+. | 2005; F06) have rotational velocities larger than $50\kms$, up to $\sim 155\kms$. |
The object studied by Shara et al. ( | The object studied by Shara et al. ( |
1997) is the second brightest BSS in 47Tuc. and all the others are located at (he low-luminosity end of the BSS region in the CMD. | 1997) is the second brightest BSS in 47Tuc, and all the others are located at the low-luminosity end of the BSS region in the CMD. |
In addition they span almost the entire range of surface (emperalures | In addition they span almost the entire range of surface temperatures |
Section 2.3)). | Section \ref{aquila:substructure}) ). |
In both cases, there is a clear correlation of stellar mass and halo mass, indicating that the processes that determine the amount of star formation per subhalo are regulated primarily by its mass. | In both cases, there is a clear correlation of stellar mass and halo mass, indicating that the processes that determine the amount of star formation per subhalo are regulated primarily by its mass. |
For a halo with an infall mass of ~10°Mo, the corresponding stellar mass is between a few times 10° to a a few times 10’Mo. | For a halo with an infall mass of $\sim10^9 \Ms$, the corresponding stellar mass is between a few times $10^5$ to a a few times $10^7 \Ms$. |
It should be noted that the minimum stellar mass resolved in the simulations is 2x10°Mo. | It should be noted that the minimum stellar mass resolved in the simulations is $2\times10^5 \Ms$. |
The two subhaloes discussed in Section 4.2,, which underwent particularly strong tidal stripping, can be identified as outliers in the relation of stellar mass to present halo mass. | The two subhaloes discussed in Section \ref{aquila:extremes}, which underwent particularly strong tidal stripping, can be identified as outliers in the relation of stellar mass to present halo mass. |
Overall, the scatter is noticeably smaller when the mass at infall, rather than the present day mass is considered, suggesting that the evolution of the satellite after infall also plays a role in some cases. | Overall, the scatter is noticeably smaller when the mass at infall, rather than the present day mass is considered, suggesting that the evolution of the satellite after infall also plays a role in some cases. |
However, it is worth noting that environmental effects primarily reduce themass, rather than themass, contrary to the scenario described in Section 1,, whereby faint dwarf galaxies are formed through stripping of baryons. | However, it is worth noting that environmental effects primarily reduce the, rather than the, contrary to the scenario described in Section \ref{introduction}, whereby faint dwarf galaxies are formed through stripping of baryons. |
Figure 8 also includes a comparison with results from our earlier simulations of isolated dwarf galaxies with much higher resolution. | Figure \ref{fig:relation_mdm-ms} also includes a comparison with results from our earlier simulations of isolated dwarf galaxies with much higher resolution. |
In both panels, the black stars denote results from simulations labeled 12-20, with total masses of 2.3x 105-10? presented in ?,, with stellar particle masses of 5.4x 10?-2.7M5,x103Mo. | In both panels, the black stars denote results from simulations labeled 12–20, with total masses of $2.3\times10^8$ $10^9\Ms$, presented in \cite{Sawala-2010}, with stellar particle masses of $5.4\times 10^2$ $2.7 \times 10^3
\Ms$. |
Blue stars are adopted from ?,, where six haloes with representative merger histories and a common mass scale of ~10?M were re-simulated, with a stellar particle mass resolution of 9x10?Mo. | Blue stars are adopted from \cite{Sawala-2011}, , where six haloes with representative merger histories and a common mass scale of $\sim10^{10}\Ms$ were re-simulated, with a stellar particle mass resolution of $9\times10^3\Ms$. |
We find that, despite the difference in resolution of up to two orders of magnitude, the results are in good agreement between the different sets of simulations, particularly when the dark matter masses are corrected for the effect of stripping, as shown in the right panel. | We find that, despite the difference in resolution of up to two orders of magnitude, the results are in good agreement between the different sets of simulations, particularly when the dark matter masses are corrected for the effect of stripping, as shown in the right panel. |
Because the same code has been used in all three sets of simulations, it follows that the results are not strongly affected by resolution. | Because the same code has been used in all three sets of simulations, it follows that the results are not strongly affected by resolution. |
The two panels in Figure 9 both show the change in stellar mass — halo mass ratio of each object from infall to the present. | The two panels in Figure \ref{fig:ratio-infall-approach} both show the change in stellar mass – halo mass ratio of each object from infall to the present. |
The ratio at infall is shown on the x-axis, while the present ratio is shown on the y-axis. | The ratio at infall is shown on the x-axis, while the present ratio is shown on the y-axis. |
Most points lie close to the black line, which indicates a constant ratio. | Most points lie close to the black line, which indicates a constant ratio. |
Notably however, the majority of haloes are above the line, meaning that their stellar mass fraction has increased since infall. | Notably however, the majority of haloes are above the line, meaning that their stellar mass fraction has increased since infall. |
This can be understood as a consequence of preferential stripping of dark matter compared to stellar matter, which is more centrally concentrated and therefore more strongly bound to the satellite. | This can be understood as a consequence of preferential stripping of dark matter compared to stellar matter, which is more centrally concentrated and therefore more strongly bound to the satellite. |
In the left panel, the colour-coding is by infall redshift; black and blue symbols indicate recent accretion, yellow and red symbols indicate infall at high redshift. | In the left panel, the colour-coding is by infall redshift; black and blue symbols indicate recent accretion, yellow and red symbols indicate infall at high redshift. |
In general, satellites that fell in earlier are more likely to have changed their ratio since infall, as expected if the change is due to continuous tidal stripping. | In general, satellites that fell in earlier are more likely to have changed their ratio since infall, as expected if the change is due to continuous tidal stripping. |
In the right panel, the colour-coding is done by distance of closest approach between the subhalo and the halo of the central galaxy. | In the right panel, the colour-coding is done by distance of closest approach between the subhalo and the halo of the central galaxy. |
As expected, haloes that had closer encounters are also the ones that underwent a slightly stronger change in the stellar mass to halo mass ratio since infall. | As expected, haloes that had closer encounters are also the ones that underwent a slightly stronger change in the stellar mass to halo mass ratio since infall. |
It appears that the haloes with the greatest distance (Dmin>300kpc) have seen no change in the ratio, but these are commonly also subhaloes that have fallen in only recently (Zing«0.2). | It appears that the haloes with the greatest distance $_{min} > 300$ kpc) have seen no change in the ratio, but these are commonly also subhaloes that have fallen in only recently $_{inf} <
0.2$ ). |
In both panels, the sizes of the symbols indicate total mass; larger satellites are typically found with higher stellar mass total mass ratios, independent of infall time or orbit. | In both panels, the sizes of the symbols indicate total mass; larger satellites are typically found with higher stellar mass -- total mass ratios, independent of infall time or orbit. |
Due to the small numbers of stellar particles per subhalo in the simulation, a detailed analysis of stellar populations is not possible. | Due to the small numbers of stellar particles per subhalo in the simulation, a detailed analysis of stellar populations is not possible. |
As a proxy for star formation history, we consider the maximum iron abundance [Fe/H] of the stars in each satellite galaxy. | As a proxy for star formation history, we consider the maximum iron abundance [Fe/H] of the stars in each satellite galaxy. |
Because iron is formed only in the late stages of stellar evolution and injected into the interstellar medium via supernovae, the amount of iron observed in stars corresponds the specific degree of reprocessing of material within each galaxy, and the intensity and duration of star formation. | Because iron is formed only in the late stages of stellar evolution and injected into the interstellar medium via supernovae, the amount of iron observed in stars corresponds the specific degree of reprocessing of material within each galaxy, and the intensity and duration of star formation. |
Figure 10. shows the maximum stellar iron abundance of the satellites, as a function of present distance (left), distance of closest approach (centre), and present-day stellar mass (right). | Figure \ref{fig:fe-relations} shows the maximum stellar iron abundance of the satellites, as a function of present distance (left), distance of closest approach (centre), and present-day stellar mass (right). |
Note that satellites with only a single generation of stars have primordial abundances, i.e. [Fe/H]=—oo, and therefore do not appear on the plotted relations. | Note that satellites with only a single generation of stars have primordial abundances, i.e. $[\mathrm{Fe/H}]\equiv- \infty$, and therefore do not appear on the plotted relations. |
The lack of a correlation on both the left and central panels indicate that the iron abundance does not depend strongly on either present distance, or distance of closest approach in the past. | The lack of a correlation on both the left and central panels indicate that the iron abundance does not depend strongly on either present distance, or distance of closest approach in the past. |
Bycontrast, there is a strong correlation with stellar mass, as observed in the Local | Bycontrast, there is a strong correlation with stellar mass, as observed in the Local |
large. distinctive ionization nebulae In nearby galaxies where optical counterparts can be identified. we may be able to detect some of the distinctive disk features of close binary SSSs (/,,5«3 days) or to detect the donor stars in wide binary SSSs (P,,5 generally is more than 100 days). ( | large, distinctive ionization nebulae In nearby galaxies where optical counterparts can be identified, we may be able to detect some of the distinctive disk features of close binary SSSs $P_{orb}< 3$ days) or to detect the donor stars in wide binary SSSs $P_{orb}$ generally is more than 100 days). ( |
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