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reftig:clr. shows the relation between the 65tun90jum. colour and the 140jum90pum colour??). | \\ref{fig:clr} shows the relation between the $65~\micron -90~\micron$ colour and the $140~\micron -90~\micron$ colour. |
The overall trend from the lower right to the upper left can be interpreted as a sequence of dust temperature. | The overall trend from the lower right to the upper left can be interpreted as a sequence of dust temperature. |
ow we theoretically quantify the observed. colour-colour relation by adopting the calculations in Hirashitaetal.(2007)... who treat the dust heating by an ISRF and the dust cooling by thermal radiation to calculate the temperature distribution function using the framework developed by Draine&Li(2001). | Now we theoretically quantify the observed colour–colour relation by adopting the calculations in \citet{hirashita07}, who treat the dust heating by an ISRF and the dust cooling by thermal radiation to calculate the temperature distribution function using the framework developed by \citet{draine01}. |
. The physical quantities that explain well the dust emission properties in the solar neighbourhood are adopted: the ISRF SED by &Panagia (1983)... the grain size distribution by Mathis.Rumpl.&Nordsieck (1977).. and the heat capacity of grain materials by Draine&Li(2001). | The physical quantities that explain well the dust emission properties in the solar neighbourhood are adopted: the ISRF SED by \citet{mathis83}, the grain size distribution by \citet{mathis77}, , and the heat capacity of grain materials by \citet{draine01}. |
. The of dust (silicate and graphite) are taken from Draine&Lee(1984) for Ax100jum and extrapolated by assuming a functional form proposed by Reach (i.e... smoothly changes from | to 2 around A~ tum). | The of dust (silicate and graphite) are taken from \citet{draine84}
for $\lambda\leq 100~\micron$ and extrapolated by assuming a functional form proposed by \citet{reach95}
$\beta$ smoothly changes from 1 to 2 around $\lambda\sim 200~\micron$ ). |
As shown by Hirashitaetal.(2007)... a slight change in 3 affects the colourcolour sequence significantly. | As shown by \citet{hirashita07}, a slight change in $\beta$ affects the colour–colour sequence significantly. |
To check the consistency with other nearby galaxies in the colour-colour diagram. adopting the same emission coefficient as adopted in Hirashitaetal.(2007) is crucial here. | To check the consistency with other nearby galaxies in the colour–colour diagram, adopting the same emission coefficient as adopted in \citet{hirashita07} is crucial here. |
We vary the ISRF with the spectral shape fixed. and denote the ISRF intensity relative to the solar neighbourhood value as 4. | We vary the ISRF with the spectral shape fixed, and denote the ISRF intensity relative to the solar neighbourhood value as $\chi$. |
We refer to Hirashitaet for the details of the framework and some basic results. | We refer to \citet{hirashita07} for the details of the framework and some basic results. |
We show the FIR colour-colour relation for various ISRF intensity x in reftig:clr.. | We show the FIR colour–colour relation for various ISRF intensity $\chi$ in \\ref{fig:clr}. |
We only show the results for graphite. since silicate follows the almost identical FIR colour-eolour relation to graphite (Hirashitaetal.2007). | We only show the results for graphite, since silicate follows the almost identical FIR colour–colour relation to graphite \citep{hirashita07}. |
. We observe that the FIR colours obtained for the individual grids are roughly explained with y=1-30. although most of the points are located systematically above the theoretical predictions on the diagram. | We observe that the FIR colours obtained for the individual grids are roughly explained with $\chi =1$ –30, although most of the points are located systematically above the theoretical predictions on the diagram. |
The lower values of x correspond to the lower-71.«; regions such as interarm regions. and the higher values to the centre and he bright spots in the spiral arms. | The lower values of $\chi$ correspond to the $T_\mathrm{LG}$ regions such as interarm regions, and the higher values to the centre and the bright spots in the spiral arms. |
The above theoretical colour-coour sequence is correct if the radiation field in a grid is approximaed to be uniform. | The above theoretical colour–colour sequence is correct if the radiation field in a grid is approximated to be uniform. |
As clearly seen in the Ha image in reffig:image.. there are small-scale star-forming regions. which should host warmer dust because of σα radiation field intensity. | As clearly seen in the $\alpha$ image in \\ref{fig:image}, there are small-scale star-forming regions, which should host warmer dust because of high radiation field intensity. |
In order to examine the effect of sucYa contamination" of warm dust. we show the FIR colour—colour relation by mixing the results with y=Land 4=100: the former value represents the general ISRF. while the latter is taken as a representative high radiation field value. | In order to examine the effect of such a `contamination' of warm dust, we show the FIR colour–colour relation by mixing the results with $\chi =1$ and $\chi =100$ : the former value represents the general ISRF, while the latter is taken as a representative high radiation field value. |
The fraction of the latter hhigher X) component is denotedas fi: that is. the intensity is calculated by where /,(x) is 1, as a function of X. | The fraction of the latter higher $\chi$ ) component is denotedas $f_\mathrm{h}$; that is, the intensity is calculated by where $I_\nu (\chi )$ is $I_\nu$ as a function of $\chi$. |
The result is shown in ο wi | The result is shown in Fig. \ref{fig:clr}. |
tte- a fraction of fi,=0.)03 significantly lift the colour sequence upwards. explaining the upper part of the FIR. colourcolour relation of the individual grids. | A slight contamination of the $\chi$ component with a fraction of $f_\mathrm{h}=0.003$ significantly lift the colour sequence upwards, explaining the upper part of the FIR colour–colour relation of the individual grids. |
This is because the 65 tum intensity responds most sensitively to the higher 4 component. | This is because the 65 $\micron$ intensity responds most sensitively to the higher $\chi$ component. |
The shift of the FIR colour-eolour relation by the contamination of a higher v component is consistent with the conclusion by Hibietal.(2006) and Hirashitaetal.(2007). | The shift of the FIR colour–colour relation by the contamination of a higher $\chi$ component is consistent with the conclusion by \citet{hibi06} and \citet{hirashita07}. |
. Hibietal.(2006) called this upper sequence “sub-correlation’. and the contamination effect "overlap effect’. | \citet{hibi06} called this upper sequence `sub-correlation', and the contamination effect `overlap effect'. |
We also compare our results with the FIR colours of the galaxies in the FIS Bright Source Catalogue in reftig:elr.. | We also compare our results with the FIR colours of the galaxies in the FIS Bright Source Catalogue in \\ref{fig:clr}. |
The analysis of these galaxies has been done by Pollo.&Takeuchi (2010). | The analysis of these galaxies has been done by \citet*{pollo10}. |
. Since there are an enormous number of galaxies. we only show the area covered by the FIR Bright Source Catalogue galaxies. | Since there are an enormous number of galaxies, we only show the area covered by the FIR Bright Source Catalogue galaxies. |
The redshifts of the sample are small and do not affect the colours. | The redshifts of the sample are small and do not affect the colours. |
We observe that the FIR colour—colour relation in 881 is within the consistent regime covered by the galaxy sample. | We observe that the FIR colour--colour relation in 81 is within the consistent regime covered by the galaxy sample. |
Note that the FIR colours of the FIS Bright Source Catalog sample present the global colours. not those in individual regions within a galaxy. | Note that the FIR colours of the FIS Bright Source Catalog sample present the global colours, not those in individual regions within a galaxy. |
Thus. we confirm that the FIR colours of individual regions within a galaxy is fundamental in determining the global galaxy colours. | Thus, we confirm that the FIR colours of individual regions within a galaxy is fundamental in determining the global galaxy colours. |
The larger scatter of the FIS Bright Source Catalogue sample may be due to a larger extent of the radiation field or a peculiarity of dust emission properties in some galaxies. | The larger scatter of the FIS Bright Source Catalogue sample may be due to a larger extent of the radiation field or a peculiarity of dust emission properties in some galaxies. |
In reffig:elr.. we also show the colours of some representative regions within the circles of |-aremin radius as shown in reffigsimage:: the central region CC). the spiral arms (S17. 7S2. and "S3Y. and the interarm regions CIF and Ἱο. | In \\ref{fig:clr}, we also show the colours of some representative regions within the circles of 1-arcmin radius as shown in \\ref{fig:image}: the central region (`C'), the spiral arms (`S1', `S2', and `S3'), and the interarm regions (`I1' and `I2'). |
We calculate he flux integrated for the circular regions. and take the flux ratios o show the colours. | We calculate the flux integrated for the circular regions, and take the flux ratios to show the colours. |
The fluxes are listed in Table [.. | The fluxes are listed in Table \ref{tab:flux}. . |
The central region have the bluest colours. while the interarm regions tend to yive redder colours than the spiral arms: these trends in colours are consistent with the temperature map shown in Fig. 5.. | The central region have the bluest colours, while the interarm regions tend to have redder colours than the spiral arms: these trends in colours are consistent with the temperature map shown in Fig.\ref{fig:dist_T_tau}. . |
Moreover. he central and interarm regions are located near to the heoretical predictions with varying ISRFs (squares in Fig. 60) | Moreover, the central and interarm regions are located near to the theoretical predictions with varying ISRFs (squares in Fig. \ref{fig:clr}) ) |
onhe colour-colour diagram. while the spiral arms are shifted toward he theoretical predictions with a mixture of general ISRF and a vigh radiation field (asterisks in Fig. 6). | onthe colour–colour diagram, while the spiral arms are shifted toward the theoretical predictions with a mixture of general ISRF and a high radiation field (asterisks in Fig. \ref{fig:clr}) ). |
This indicates that the | This indicates that the |
tomography allows for a straightforward investigation of the redshift-dependence of the shear signal. | tomography allows for a straightforward investigation of the redshift-dependence of the shear signal. |
Consequently. we conclude that both approaches have their advantages ancl it is sensible to do both. | Consequently we conclude that both approaches have their advantages and it is sensible to do both. |
We especially thank Rongmon Bordoloi or use of the photometric redshift estimates from Borcolot et al.. ( | We especially thank Rongmon Bordoloi for use of the photometric redshift estimates from Bordoloi et al., ( |
2009) and Filipe Xbdalla for use of redshift estimates rom Abdalla et al. ( | 2009) and Filipe Abdalla for use of redshift estimates from Abdalla et al., ( |
2007) in an earlier craft. | 2007) in an earlier draft. |
We thank Adam Amara for a careful reading of an carly draft. and an anonymous referee for many helpful comments. | We thank Adam Amara for a careful reading of an early draft and an anonymous referee for many helpful comments. |
We also hank Andy Taylor anc Fergus Simpson for many useful discussions. | We also thank Andy Taylor and Fergus Simpson for many useful discussions. |
TDIx is supported by a STEC Rolling Grant LAOSSS. | TDK is supported by a STFC Rolling Grant RA0888. |
li is well known that the masses of the supermassive black holes (ΛΙΠΗΣΕΞ} in the nuclei of carlytype galaxies and bulges correlate with the velocity dispersions of the stellar spheroids (e.g..Gebhardtetal.2000:Ferrarese&|Merritt2000:Tremaineetal. 2002). | It is well known that the masses of the supermassive black holes (SMBHs) in the nuclei of early–type galaxies and bulges correlate with the velocity dispersions of the stellar spheroids \citep[e.g.,][]{Gebhardt00, Ferrarese00,
Tremaine02}. |
. A simple explanation invokes momentum feedback (Ixing2003.2005). | A simple explanation invokes momentum feedback \citep{King03,King05}. |
. In this picture the SALDAIL luminosity is limited. by the Ecldineton value. and 1o momentum outflow rate produced by radiation pressure is of the order of where & is the opacity. assumed. to be dominated: by the electron scattering. and. Mpg is the SMDII mass. | In this picture the SMBH luminosity is limited by the Eddington value, and the momentum outflow rate produced by radiation pressure is of the order of where $\kappa$ is the opacity, assumed to be dominated by the electron scattering, and $M_{\rm BH}$ is the SMBH mass. |
This momentum flux produces an outward force on the gas in the bulge. whose weight is WOR)=CALGR)Ανα]IU. where MGR) is the enclosed. gas. mass at radius 2. and Missa) is the total enclosed mass including dark matter. | This momentum flux produces an outward force on the gas in the bulge, whose weight is $W(R) = GM(R)[M_{\rm total}(R)]/ R^2$, where $M(R)$ is the enclosed gas mass at radius $R$, and $M_{\rm total}(R)$ is the total enclosed mass including dark matter. |
For an isothermal potential ACR) and Ata?) are proportional to Εν so the result is Llere f, is the barvonic fraction. aud o=GMÍ2R is the velocity dispersion in the bulge. | For an isothermal potential, $M(R)$ and $M_{\rm total}(R)$ are proportional to $R$, so the result is Here $f_g$ is the baryonic fraction and $\sigma^2 = GM/2R$ is the velocity dispersion in the bulge. |
To order of magnitude. the relation 2. holds for any potential if estimated. at the virial raclius. | To order of magnitude, the relation \ref{w} holds for any potential if estimated at the virial radius. |
Requiring that momentum output produced. by. the black hole should just balance the weight of the gas Leads to the Mpg (0 relation (Ixing2003.2005):: The model is attractive. in its physical simplicity. | Requiring that momentum output produced by the black hole should just balance the weight of the gas leads to the $M_{\rm BH}$ $\sigma$ relation \citep{King03,King05}: The model is attractive in its physical simplicity. |
Further. the result contains no free parameters. but is very close to the observed. Adpy σ relation. | Further, the result contains no free parameters, but is very close to the observed $M_{\rm BH}$ $\sigma$ relation. |
Another feature commonly found in. the. centres of ealaxies are nuclear star clusters. | Another feature commonly found in the centres of galaxies are nuclear star clusters. |
They are found in late type spirals (e.g..Bokeretal. 2002).. bulgeless spirals | They are found in late type spirals \citep[e.g.,][]{BokerEtal02}, , bulgeless spirals |
Iu Brooketal.(200[b).. we found that the thick disk scale-Ieugth of our simulated galaxy was srorter than that of the thin disk. | In \citet{brook04b}, we found that the thick disk scale-length of our simulated galaxy was shorter than that of the thin disk. |
The observations ο" Yoachimi&Daleanton(2006) fiud that thick disk scale-lcsneths are svsteinaticallv larger than those of thin «isks. | The observations of \cite{yoach2} find that thick disk scale-lengths are systematically larger than those of thin disks. |
Iu this meregcr simulation. the hot merger star poplation has an exponential profile with scale-Ieugth larger than that of the later-formine disk star population. | In this merger simulation, the hot merger star population has an exponential profile with scale-length larger than that of the later-forming disk star population. |
This lay favor a significantoO oOeas-rich mereerOo beineC» a feature of disk Oogalaxw formation. | This may favor a significant gas-rich merger being a feature of disk galaxy formation. |
Furt∐∖↥⋅↴∖↴↕⋯∏↕⋜↧↑↕∪ is will detenine whether old. hot disks with scale-leneth larecr than those of voung cold disss result for a wide range of nierecr parameters. | Further simulations will determine whether old, hot disks with scale-length larger than those of young cold disks result for a wide range of merger parameters. |
Caution is required in interpreting the result. as subsequeit “inside-out” thin lish erowth through eas infall (uot present iu the current simulation) may increase the thin disk scale-lenetl (Breοκetal.2006b). | Caution is required in interpreting the result, as subsequent “inside-out” thin disk growth through gas infall (not present in the current simulation) may increase the thin disk scale-length \citep{brook06b}. |
. Yoachiu&Dalcautou(2006) further fud that low-mass disk οslaxies have larger thiczthin disk mass ratios. | \cite{yoach2} further find that low-mass disk galaxies have larger thick:thin disk mass ratios. |
They interpret this as evidence for the formation of the hick disk by clirect accrοποια of stars. as progenitors of lower mass galaxies will more casily expel their οeas from their simall potential well prior to the merger. | They interpret this as evidence for the formation of the thick disk by direct accretion of stars, as progenitors of lower mass galaxies will more easily expel their gas from their small potential well prior to the merger. |
Yet observations suggest that low mass galaxies are in fact nore eas-rich. both at low redshift(οι, Schombert.MeCGaughl.&Eder 2001)) and high redshift (Erbeal.2006)). favoring eas-rich merecrs as an interpretation of the higher thick:thin disk mass ratio iu low mass galaxies. | Yet observations suggest that low mass galaxies are in fact more gas-rich, both at low redshift (e.g. \citealt{schombert}) ) and high redshift \citealt{erb}) ), favoring gas-rich mergers as an interpretation of the higher thick:thin disk mass ratio in low mass galaxies. |
Also. the erowth of the thin disk in suall galaxies is perhaps regulated by their low deusitics (Dalcanton2006)) alc supernovae feedback. 1.0. the high thick to thin disk ratio of low-mass galaxies may be the result of less erowthn of a thin disk. | Also, the growth of the thin disk in small galaxies is perhaps regulated by their low densities \citealt{dalcanton06}) ) and supernovae feedback, i.e. the high thick to thin disk ratio of low-mass galaxies may be the result of less growth of a thin disk. |
- Our idealized study does not include cold accretion from the intergalactic medimu aud infall processes. which are miportaut in the birth of disk galaxies. | Our idealized study does not include cold accretion from the intergalactic medium and infall processes, which are important in the birth of disk galaxies. |
Further. iu the hierarchical structure formation scenario. a range of merecr histories exists. aud to explain the predominance of thick disk componeuts in disk ealaxies. One cannot rely on a sinele merecr eveut. | Further, in the hierarchical structure formation scenario, a range of merger histories exists, and to explain the predominance of thick disk components in disk galaxies, one cannot rely on a single merger event. |
Rather. à few or several siguificant eas-rich merecr events are likely to occur at an carly epoch in the formation of a disk galaxy (Brooketal.2005 ). | Rather, a few or several significant gas-rich merger events are likely to occur at an early epoch in the formation of a disk galaxy \citealt{brook05}) ). |
A cold mode of accrοποια from flamentary structures also occurs in this scenario. | A cold mode of accretion from filamentary structures also occurs in this scenario. |
But our study supports the view that the violent accretion of Ooeas-ricli Oogalaxies is ceutral iu producingC» thick disk properties. | But our study supports the view that the violent accretion of gas-rich galaxies is central in producing thick disk properties. |
Dy ieunoriigOoC» cold accretion. or siupli&ed study hiehliehts the effect of the starburst associated with such mergers on the chemical abuudances of the foriuug stars. iu particular the hieh e clement abundances at high metallicities. aud vertical abundance gracicuts. | By ignoring cold accretion, our simplified study highlights the effect of the starburst associated with such mergers on the chemical abundances of the forming stars, in particular the high $\alpha$ element abundances at high metallicities, and vertical abundance gradients. |
It is vet to be shown that cold accretion alouc. perhaps resulting in a thin disk which is heated by au iufalling satelli sor the dispersion of large star clusters (Ikroupa 2006)). can reproduce such chemical signatures. | It is yet to be shown that cold accretion alone, perhaps resulting in a thin disk which is heated by an infalling satellite, or the dispersion of large star clusters \citealt{kroupa,elmegreen}) ), can reproduce such chemical signatures. |
Cas-rich iiergers as the dominant process in foriiug thick disks provide a natural explanation of observed abundance patterns and egradieuts iu the Milky Way disk components. | Gas-rich mergers as the dominant process in forming thick disks provide a natural explanation of observed abundance patterns and gradients in the Milky Way disk components. |
A high-redshift epoch of eas-ricli mereers is emiereiug as an important phase in the VOYV birth of disk galaxies. | A high-redshift epoch of gas-rich mergers is emerging as an important phase in the very birth of disk galaxies. |
The simulation was performed at the Laboratoire d'Astroplivsique Nunérrique. Université Laval. | The simulation was performed at the Laboratoire d'Astrophysique Numérrique, Université Laval. |
CD. SR. TAL are funded by the Canada Research Chair program aud NSERC. | CB, SR, HM are funded by the Canada Research Chair program and NSERC. |
DIN is a JSPS Fellow. | DK is a JSPS Fellow. |
obtained taking into account galaxies in clusters at ς<0.1 and brighter than M,=—20. and those located 1n clusters atz«0.07 and brighter than M,=—19. | obtained taking into account galaxies in clusters at $z<0.1$ and brighter than $_{r}=-20$, and those located in clusters at $z<0.07$ and brighter than $_{r}=-19$. |
No differences in the clusters properties were observed in the two samples of galaxies (see Tab. | No differences in the clusters properties were observed in the two samples of galaxies (see Tab. |
2). | 2). |
Figure 2 shows the cumulative distribution functions of «c. fp and Aun» for cluster with and without substructure. | Figure \ref{subes} shows the cumulative distribution functions of $\sigma_{c}$, $f_{b}$ and $\Delta m_{12}$ for cluster with and without substructure. |
This figure also shows the cumulative distribution functions of the number of galaxy members in clusters with and without substructure. | This figure also shows the cumulative distribution functions of the number of galaxy members in clusters with and without substructure. |
Notice that the number of clusters with substructure is very small when galaxies brighter than M,=-20.0 are considered. | Notice that the number of clusters with substructure is very small when galaxies brighter than $M_{r}=-20.0$ are considered. |
However. this number is larger when we study galaxies brighter than M,= —]9.0. | However, this number is larger when we study galaxies brighter than $M_{r}=-19.0$ . |
In this case. the Kolmogorov-Smirnoff (KS) test does not report statistical differences in the cluster properties for clusters with and without substructure. | In this case, the Kolmogorov-Smirnoff (KS) test does not report statistical differences in the cluster properties for clusters with and without substructure. |
Thus. substructure in the inner cluster region does not affect οι. fj or Aun». | Thus, substructure in the inner cluster region does not affect $\sigma_{c}$, $f_{b}$ or $\Delta m_{12}$. |
The above result is not in agreement with Ramella et al. ( | The above result is not in agreement with Ramella et al. ( |
2007). | 2007). |
They found a clear difference between the mean value of Ας for cltsters with and without substructure. | They found a clear difference between the mean value of $\Delta m_{12}$ for clusters with and without substructure. |
The work by Ramella et al. | The work by Ramella et al. |
was based on a sample of 77 nearby clusters (O.04<= <0.07) from the WINGS survey (Fasano et al. | was based on a sample of 77 nearby clusters $<z<$ 0.07) from the WINGS survey (Fasano et al. |
2006). | 2006). |
The substrecture was determined using the DEDICA procedure (Pisani 1993. 1996). | The substructure was determined using the DEDICA procedure (Pisani 1993, 1996). |
This procedure was aplied to all galaxies in the cluster sample brighter than My=—16. | This procedure was aplied to all galaxies in the cluster sample brighter than $M_{V}=-16$. |
The different galaxy population studied could be the reason of the disagreement with Ramella et al. ( | The different galaxy population studied could be the reason of the disagreement with Ramella et al. ( |
2007). | 2007). |
Instead of summing all individual values of 6 to get an estimate of the amount of substructure in the clusters as a whole. we can use the values of 9 for each galaxy in order to select those galaxies located in substructures. | Instead of summing all individual values of $\delta$ to get an estimate of the amount of substructure in the clusters as a whole, we can use the values of $\delta$ for each galaxy in order to select those galaxies located in substructures. |
The problem is to determine the value 9,. which effectively separates galaxies inside and outside substructures. | The problem is to determine the value $\delta_{c}$, which effectively separates galaxies inside and outside substructures. |
We have taken into account two different approaches to this problem. | We have taken into account two different approaches to this problem. |
The first approach is to consider different values of 0; for each cluster. | The first approach is to consider different values of $\delta_{i}$ for each cluster. |
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