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Optically Chick dust emission in ULIRGs at ΤαΗν wavelengths has been first noticed by Condon οἱ al. | Optically thick dust emission in ULIRGs at far-IR wavelengths has been first noticed by Condon et al. |
1991 using interferometric radio continuum imaging to identify the true sizes of (heir starburst regions. | 1991 using interferometric radio continuum imaging to identify the true sizes of their starburst regions. |
To explore its effect on emergent dust emission we follow Lisenleld et al. | To explore its effect on emergent dust emission we follow Lisenfeld et al. |
2000 ancl replace (he SE-heated and cold dust components inside the brackets of Equation 7 with where τν)= r(r/v.y. | 2000 and replace the SF-heated and cold dust components inside the brackets of Equation 7 with where $\rm \tau (\nu ) = \tau _{\circ} (\nu/\nu_{\circ})^2$. |
lIn Table 3 we list the results of these dust SED fits. and in Figure 6 show the distribution of the deduced optical cepts for all (he LIBGs with available submm data from the literature. | In Table 3 we list the results of these dust SED fits, and in Figure 6 show the distribution of the deduced optical depths for all the LIRGs with available submm data from the literature. |
Substantial dust optical depths are found at 100 (~4 for most LIRGs. in accord with earlier results (e.g. Solomon et al. | Substantial dust optical depths are found at $\mu $ m $\sim $ 4--21) for most LIRGs, in accord with earlier results (e.g. Solomon et al. |
1997: Lisenleld et al. | 1997; Lisenfeld et al. |
2000). | 2000). |
Even 11063. a vigorously star-forming LIRG where a massive and extended cold dust reservoir wilhTy~(10 IWIN is revealed via subi imagine (Papadopoulos Seaquist 1999). can have its non-AGN global dust emission fitted equally well with a single-temperature but oplically thick SED. | Even 1068, a vigorously star-forming LIRG where a massive and extended cold dust reservoir $\rm
T_{d}$$\sim$ K is revealed via submm imaging (Papadopoulos Seaquist 1999), can have its non-AGN global dust emission fitted equally well with a single-temperature but optically thick SED. |
The latter would then erroneously interpret its cold dust as optically thick emission. with 7j0045774.6. | The latter would then erroneously interpret its cold dust as optically thick emission with $\rm \tau _{100\mu m}$$\sim $ 4.6. |
Thus the SED fits using Equations 7 | Thus the SED fits using Equations 7 |
For two planets to mutually excite their eccentricdiües by resonant interaction. the duration of passage through the resonance must be longer (han (he resonant libration period (see. e.g.. Dermott. Malhotra. Murray. 1933): where Ad... is the width of the resonance and 7; is the resonant libration period. | For two planets to mutually excite their eccentricities by resonant interaction, the duration of passage through the resonance must be longer than the resonant libration period (see, e.g., Dermott, Malhotra, Murray 1988): where $\Delta a_{\rm res}$ is the width of the resonance and $T_l$ is the resonant libration period. |
This requirement is tested numerically in refnumex.. | This requirement is tested numerically in \\ref{numex}. |
With no important loss of generality. let us consider the migration of the inner planet through a resonance of order |g] and take the outer planet to be the perturber on a fixed orbit. | With no important loss of generality, let us consider the migration of the inner planet through a resonance of order $|q|$ and take the outer planet to be the perturber on a fixed orbit. |
To order of magnitude. These expressions are appropriate for οSS0.8 orbits in |q|>1 resonances. and 0.32€=0.16U5/M,)7e0.01CMSMj)P orbits in the 2:1 resonance [Dermott et al. | To order of magnitude, These expressions are appropriate for $e \lesssim 0.3$ orbits in $|q| > 1$ resonances, and $0.3 \gtrsim e \gtrsim 0.1
(M_2/M_{\ast})^{1/3} \approx 0.01 (M_2/M_J)^{1/3}$ orbits in the 2:1 resonance [Dermott et al. |
1088: Murray Dermott 1999: see the latter's equations (8.58) and (8.76)]. | 1988; Murray Dermott 1999; see the latter's equations (8.58) and (8.76)]. |
Insertion of (5)) and (6)) into (4)) vields Now it is usually remarked (see. e.g.. Ward 1997 and GS) that for Type IH drift. /v. the viscous diffusion time of the disk. | Insertion of \ref{tl}) ) and \ref{dares}) ) into \ref{slow}) ) yields Now it is usually remarked (see, e.g., Ward 1997 and GS) that for Type II drift, $|a/\dot{a}| \sim a^2 / \nu$ , the viscous diffusion time of the disk. |
But this statement cannot be (rue in the limit that the planet mass greatly exceeds (he ring mass. | But this statement cannot be true in the limit that the planet mass greatly exceeds the ring mass. |
We generalize the Type II drift velocity by setting the viscous torque equal to the rate of change of angular momentum of either the planet or (he ring. whichever is more massive: | We generalize the Type II drift velocity by setting the viscous torque equal to the rate of change of angular momentum of either the planet or the ring, whichever is more massive: |
using the interpolated surface. | using the interpolated surface. |
The rMSFE method was be applied to our set of isochrones. with different a Ile content and/or heavy element distribution. bv shifting in each case (he considered isochrone in both color and magnitude to fit the corresponding same metallicityreference one. | The rMSF method was be applied to our set of isochrones, with different a He content and/or heavy element distribution, by shifting in each case the considered isochrone in both color and magnitude to fit the corresponding same metallicity one. |
This procedure provides the brightness of the ISTO relative to thereference isochrones MSTO. which can be used to derive the relative age as previously described. | This procedure provides the brightness of the MSTO relative to the isochrone's MSTO, which can be used to derive the relative age as previously described. |
The impact of an enhanced initial He content or an extreme CNONa mixture on the relative ages obtained by the different dating methocls is illustrated in Figure3. | The impact of an enhanced initial He content or an extreme CNONa mixture on the relative ages obtained by the different dating methods is illustrated in Figure. |
Horizontal. vertical. and YMSE. methods applied to isochrones with the same |Fe/1l]. same age. but different. values of Y (upper panels) and different. CNONa abundances. (lower panel) are shown. | Horizontal, vertical, and rMSF methods applied to isochrones with the same [Fe/H], same age, but different values of Y (upper panels) and different CNONa abundances (lower panel) are shown. |
The upper left hand panel shows how. for the same age and same [Fe/Il]. the horizontal parameter increases when decreasing This means that the horizontal-method derived age does not coincide with the isochrone (input) one. | The upper left hand panel shows how, for the same age and same [Fe/H], the horizontal parameter increases when decreasing This means that the horizontal-method derived age does not coincide with the isochrone (input) one. |
1t depends on the Y value. the derived age being older for larger values of Y. The upper central panel illustrates the effect ol Y on the vertical method. | It depends on the Y value, the derived age being older for larger values of Y. The upper central panel illustrates the effect of Y on the vertical method. |
The large impact that the He content has on the ZAIID level. as well as on the MSTO brightness. can be seen. | The large impact that the He content has on the ZAHB level, as well as on the MSTO brightness, can be seen. |
This translates into a huge effect on the vertical parameter and consequently on the derived ages. | This translates into a huge effect on the vertical parameter and consequently on the derived ages. |
It is worth mentioning that the ZAIID and MSTO magnitude shifts as a consequence of Y variation do not compensate each other. | It is worth mentioning that the ZAHB and MSTO magnitude shifts as a consequence of Y variation do not compensate each other. |
On the contrary. increasing Y translates into a fainter MSTO and a brighter ZAIILD. | On the contrary, increasing Y translates into a fainter MSTO and a brighter ZAHB. |
The effect of varving the He abundance on the rMSE method is illustrated in the upper right hand panel. | The effect of varying the He abundance on the rMSF method is illustrated in the upper right hand panel. |
Increasing Y makes the MSTO [ainter. as well as the AIS locus. | Increasing Y makes the MSTO fainter, as well as the MS locus. |
According to the figure. both effects compensate each other during the rMSE. procedure. so the method (urns out to be rather insensitive to Πο variations. | According to the figure, both effects compensate each other during the rMSF procedure, so the method turns out to be rather insensitive to He variations. |
Similar arguments can be made for the CNONa variations based on the lower panels. | Similar arguments can be made for the CNONa variations based on the lower panels. |
The lower left hand panel illustrates the effect that a different CNONa abundance has on the horizontal-method derived age. | The lower left hand panel illustrates the effect that a different CNONa abundance has on the horizontal-method derived age. |
In the example of the figure. the horizontal parameter | In the example of the figure, the horizontal parameter |
οςτα could be plausible. | medium could be plausible. |
11291. hosts many ieelons. a quiescent nucleus. and two rather faint floceuleut spiral indus. | 4294 hosts many regions, a quiescent nucleus, and two rather faint flocculent spiral arms. |
The outer regions of the spiral structure are consistent with cuhanced star formation at the radius at which the eas surface density iu the disk beconies πηραοΊσα for star formation (?7?).. | The outer regions of the spiral structure are consistent with enhanced star formation at the radius at which the gas surface density in the disk becomes super-critical for star formation \citep[][]{ThornleyWilson1995, Thornley1996}. |
The observed vvolocities show streaming motions aloug the spiral arms CU.andthiswork).. | The observed velocities show streaming motions along the spiral arms \citep[][ and this work]{Cheminetal2006}. |
Although the receding part of the disk appears sheltly perturbed. the overall appearance of the velocity field is quite regular. | Although the receding part of the disk appears slightly perturbed, the overall appearance of the velocity field is quite regular. |
The rotation curve is consistent with that preseuted by ?.. | The rotation curve is consistent with that presented by \citet{Rubinetal1999}. |
Usine the nuncthod. we derive O,=133,52 with MCR)x1.9 kpe in agreement with the location where the spiral arius cieree from the ends of the har. | Using the method, we derive $\Omega_p = 43_{-12}^{+3}$, with $r(CR) \approx 1.9$ kpc in agreement with the location where the spiral arms emerge from the ends of the bar. |
Within the (CR). the velocity cispersion follows a plateau. whereas outside this radius it decreases hosts many iregious iu the bar as well as the two partly broken spiral arns. | Within the $r(CR)$, the velocity dispersion follows a plateau, whereas outside this radius it decreases hosts many regions in the bar as well as the two partly broken spiral arms. |
The ceutral regions exhibit significant molecular eas content (?).. although the star formation efficiency is constant iu the inner two effective radii (2).. | The central regions exhibit significant molecular gas content \citep{Bokeretal2002}, although the star formation efficiency is constant in the inner two effective radii \citep{RowndYoung1999}. |
Our velocity field displavs a "S-shaped zero-velocity curve over the central L kpe radius region of the galaxy. | Our velocity field displays a “S”-shaped zero-velocity curve over the central 4 kpc radius region of the galaxy. |
The rotation curve rises steeply in the ceutral few arcseconds. which indicates a rapidly rotating component within =30"., and it contiuues to increase to 130 aat 6 kpe radius. | The rotation curve rises steeply in the central few arcseconds, which indicates a rapidly rotating component within $\approx 30$, and it continues to increase to 130 at 6 kpc radius. |
This is consistent with the optical rotation curve from ? aud the ppositiou-velocity diagram of ον, | This is consistent with the optical rotation curve from \citet{Rubinetal1999} and the position-velocity diagram of \citet{Helouetal1984}. |
We derive ο=20BESTES L.sehieh we assign to outer spiral aris. | We derive $\Omega_p = 20_{-8}^{+6}$, which we assign to outer spiral arms. |
We use this tto locate an outer ILR radius at Q,,about 1.5 κρο (207)) and an inner ILR radius at around 200 pc. | We use this to locate an outer ILR radius at about 1.5 kpc ) and an inner ILR radius at around 200 pc. |
The outer ILR radius is. to within the errors. consistent with the dadisk scale leusth derived by ο | The outer ILR radius is, to within the errors, consistent with the disk scale length derived by \citet{Koopmannetal2006}. |
Sinuilar to ον, we apply the uuuethod on the pixels iuterior to the ILR. aud find thatthe iuner bar is decoupled from the outer oval. as its pattern speed is Q,=15+. | Similar to \citet{Fathietal2007TW}, we apply the method on the pixels interior to the ILR, and find thatthe inner bar is decoupled from the outer oval, as its pattern speed is $\Omega_p = 45$. |
This yvields that the corotation of the immer bar is locatedQ, at ILR of the outer oval. | This yields that the corotation of the inner bar is located at ILR of the outer oval. |
The method we have applied here enables us to doteruiue the secondary. oonlv with au nuucertaiutv of504.. ©,since to zero-th order. the nunethod applied to the ceutral region has to assume that the outer bar or spiral arius are stationary. | The method we have applied here enables us to determine the secondary only with an uncertainty of, since to zero-th order, the method applied to the central region has to assume that the outer bar or spiral arms are stationary. |
This is not a realistic scenario. so we use this higher pattern speed for the inner bar only as an indication that we can detect the phenomenon. aud that the patter speed is higher than that of the outer is a graud design. possibly post starburst. spiral ealaxy with a LINER nucleus (???).. | This is not a realistic scenario, so we use this higher pattern speed for the inner bar only as an indication that we can detect the phenomenon, and that the pattern speed is higher than that of the outer is a grand design, possibly post starburst, spiral galaxy with a LINER nucleus \citep{Rushetal1993, Koornneef1993, Elfhagetal1996}. |
It hosts a prominent 2 kpc stellar bar =15 deerees from the orientation of the outer disk iajor-axis with no sienificant change when comparing D aud Jf-hand nuages (?).. and displavs a star formation rate of 0.9 citepConzalezDelgadol997.. | It hosts a prominent 2 kpc stellar bar $\approx 45$ degrees from the orientation of the outer disk major-axis with no significant change when comparing $B$ and $H$ -band images \citep{Eskridgeetal2002}, and displays a star formation rate of $0.9$ \\citep{GonzalezDelgado1997}. . |
The rotation curve was measured by? from which they derived au exponcutial main disk scale leneth of ©8.7 kpc. aud a massive bulec (Mutκκ). | The rotation curve was measured by \citet{ZasovSilchenko1987} from which they derived an exponential main disk scale length of $\approx 8.7$ kpc, and a massive bulge $M_{bulge} \geq M_{disk}$ ). |
Although the outer parts of their rotation curve are not well determined. the ner regious agree well with our rrotation curvo. | Although the outer parts of their rotation curve are not well determined, the inner regions agree well with our rotation curve. |
Our observations display a deficiency. of eenission across the bar. iu agreement with the iniges from ?7.. | Our observations display a deficiency of emission across the bar, in agreement with the images from \citet{GonzalezDelgadoPerez1997}. |
The velocity field displavs clear disk-like rotation. with various “wigeles” im the zero-volocitv curve. often interpreted as streaming motions along spiral axius (?).. | The velocity field displays clear disk-like rotation, with various “wiggles” in the zero-velocity curve, often interpreted as streaming motions along spiral arms \citep{Emsellemetal2006}. |
The rotation curve rises steadily in the central307. and reaches 250 aat 15 kpe radius, | The rotation curve rises steadily in the central, and reaches 250 at 15 kpc radius. |
We derive O,=1152 παλιά the r(CR)z17 kpe. ie.937. | We derive $\Omega_p= 14_{-1}^{+5}$ and the $r(CR) \approx 17$ kpc, i.e.,. |
. The deficiency of the ecluission from the stellay bar. combined with its optimal oricutation compared with the outer disk axis. indicate that we are probing the oof the spiral arms. aud not the well-known 2 kpe bur. | The deficiency of the emission from the stellar bar, combined with its non-optimal orientation compared with the outer disk major-axis, indicate that we are probing the of the spiral arms, and not the well-known 2 kpc bar. |
We note that bv placing corotation at 17 kpc. we bring the inner Ll resonance close to the end of the svuunetre part of the spiral structure. | We note that by placing corotation at 17 kpc, we bring the inner 4:1 resonance close to the end of the symmetric part of the spiral structure. |
At this region. in agreenient with the predictious by ?.. à bifurcation of the northern aria is seen in Fig. l.. | At this region, in agreement with the predictions by \citet{Patsisetal1997}, a bifurcation of the northern arm is seen in Fig. \ref{fig:allmaps}. |
Furthermore. the association of the end of the svuuuetric part of the spirals with the ner Ll resonance is supported by the radial variation of the ratio of the amplitudes of the m=| to the wm=2 componcuts (??).. | Furthermore, the association of the end of the symmetric part of the spirals with the inner 4:1 resonance is supported by the radial variation of the ratio of the amplitudes of the $m=4$ to the $m=2$ components \citep{Grosbol1985, Patsis1991}. |
has a very poor molecular aud neutral gas content (2). and relatively poor jonised eas ina zc812 kpc prominent stellar bar which is enclosed by sharp inner rine-like structure. | has a very poor molecular and neutral gas content \citep{Verter1985} and relatively poor ionised gas in a $\approx 8-12$ kpc prominent stellar bar which is enclosed by sharp inner ring-like structure. |
The bar coutains two straight dust lanes. and the spiral arms host many iregious (?777).. and form the rime at the bar ends. just before they become verv open. broad. and diffuse. | The bar contains two straight dust lanes, and the spiral arms host many regions \citep{GonzalezDelgadoPerez1997, Aguerrietal1998, Rautiainenetal2005}, and form the ring at the bar ends, just before they become very open, broad, and diffuse. |
At the north eud of the bar. the arma becomes brighter as it leaves the rime (?).. | At the north end of the bar, the arm becomes brighter as it leaves the ring \citep{Eskridgeetal2002}. |
7. found that the stellar bar has au axis ratio of 0.5 and is oriented. 18 degrees in the nortli- direction. | \citet{MartinFriedli1997} found that the stellar bar has an axis ratio of 0.5 and is oriented $18$ degrees in the north-east direction. |
These authors also derived au asviinetric star forination efficiency 0,2 Hn the bar. and seven times higher value in the cireunmnmclear region (?).. | These authors also derived an asymmetric star formation efficiency $<0.2$ in the bar, and seven times higher value in the circumnuclear region \citep{GonzalezDelgado1997}. . |
Like iu 55371. the eenüsson Is so weak in the bar of this galaxw that we have to make use of large spatial bius to measure aly kinematic information. | Like in 5371, the emission is so weak in the bar of this galaxy that we have to make use of large spatial bins to measure any kinematic information. |
The rotation curve caunot be constrained within the bar. | The rotation curve cannot be constrained within the bar. |
The nuuethod thus delivers the pattern speed of the prominent spiral structure (ο=13E 44). with the spiral pattern r(CR)τί8 pe (or65".ingood ?).. which is just outsidethe iradiusof the strong stellar bar (?).. | The method thus delivers the pattern speed of the prominent spiral structure $\Omega_p= 13_{-2}^{+2}$ ), with the spiral pattern $r(CR)\approx 8$ kpc \citep[or 65\arcsec, in good agreement with the $r(CR)=71$\arcsec\ of][]{RSL2008}, , which is just outsidethe radiusof the strong stellar bar \citep{ButaZhang2009}. |
This suggests that the bar shares the same pattern speed as the spiral avis (e.g.T). | This suggests that the bar shares the same pattern speed as the spiral arms \citep[e.g., ][]{Taggeretal1987}. |
hosts a very prominent aud elongated bar surrounded by diffuse. exteuded. and fragmented spiral arms with no regular pattern (2).. | hosts a very prominent and elongated bar surrounded by diffuse, extended, and fragmented spiral arms with no regular pattern \citep{ElmegreenElmegreen1987}. . |
This fiocculeut galaxy | This flocculent galaxy |
We should add a caution here. which is that for the clusters deseribed i Table l| aud plotted in Figure 2.. we have computed rey. ὃν and 7g Using the full ionizing DIuuinositv of the cluster. | We should add a caution here, which is that for the clusters described in Table \ref{clusterlist} and plotted in Figure \ref{clustersample}, we have computed $r_{\rm ch}$ , $\zeta$, and $r_{\rm stall}$ using the full ionizing luminosity of the cluster. |
This is not correct very carly in the expansion process. since the initial Strouuneren radius around each massive star is so sniall that it may enclose at most a few of its neighbors. | This is not correct very early in the expansion process, since the initial Strömmgren radius around each massive star is so small that it may enclose at most a few of its neighbors. |
Since this will lower the ionizing huninosity compared to our value. the very carly expansion could be gas-driveu. | Since this will lower the ionizing luminosity compared to our value, the very early expansion could be gas-driven. |
However. as the expanding shells overlap. the incorporation of more stars iuto their interiors will rapidly convert the expansion to a radiatively-driven one: since ςκSt. the expaudius shell in a cluster for whichwe have conrputed ¢=LOO reaches à21 aud becomes radiatiou-dominated when the shell includes onlv L6% of the cluster luminosity. | However, as the expanding shells overlap, the incorporation of more stars into their interiors will rapidly convert the expansion to a radiatively-driven one; since $\zeta \propto S_{49}^{2/3}$, the expanding shell in a cluster for whichwe have computed $\zeta = 100$ reaches $\zeta > 1$ and becomes radiation-dominated when the shell includes only $4.6\%$ of the cluster luminosity. |
Since Iuninositv teuds to be stronelv ceutrally couceutrated. this will occur well before the shell iucludes this fraction of the cloud volume. | Since luminosity tends to be strongly centrally concentrated, this will occur well before the shell includes this fraction of the cloud volume. |
Thus radiation takes over very carly in the expansion process for those clusters for which we have computed values of C» | Thus radiation takes over very early in the expansion process for those clusters for which we have computed values of $\zeta \gg 1$. |
To study the cwnamics of an expaucding region with significant raciation pressure. it is convenient to uon-dinensioualize theequation of motion (1)). | To study the dynamics of an expanding region with significant radiation pressure, it is convenient to non-dimensionalize theequation of motion \ref{momeqn}) ). |
We let =Pray and 7= t/ta,. where | We let $x=r/r_{\rm ch}$ and $\tau=t/t_{\rm ch}$ , where |
The strong correlations between the masses of central black holes (19115) and the luminosities. dynamical masses. anc velocity dispersions σ of their host galaxies imply that the growth. processes of BUs and their hosts are. intimately linked. (e. Magorrian et al. | The strong correlations between the masses of central black holes (BHs) and the luminosities, dynamical masses, and velocity dispersions $\sigma$ of their host galaxies imply that the growth processes of BHs and their hosts are intimately linked (e.g., Magorrian et al. |
1998: Ferrarese Merrit 2000: Gebhardt ct 22000: Ferrarese 2002: Ferrarese Lord 2005: Crab 2007: Tundo et al. | 1998; Ferrarese Merritt 2000; Gebhardt et 2000; Ferrarese 2002; Ferrarese Ford 2005; Graham 2007; Tundo et al. |
2007: Shankar οἱ al. | 2007; Shankar et al. |
2009b). | 2009b). |
LLowever. constraining the cosmological evolution of BUs remains a challenge. | However, constraining the cosmological evolution of BHs remains a challenge. |
Although a variety of theoretica models may roughly match observations. the underlving physical assumptions on DII growth. can vary. drastically from one model to another (e.g. Soltan 1982: Silk Rees 1998: Salucci et al. | Although a variety of theoretical models may roughly match observations, the underlying physical assumptions on BH growth can vary drastically from one model to another (e.g., tan 1982; Silk Rees 1998; Salucci et al. |
1999: Cavaliere Vittorini 2000: Kaullmann Llaehnelt 2000; Yu ‘Tremaine 2002: Steed Weinberg 2003: Wyithe Loch 003: Ciranato et al. | 1999; Cavaliere Vittorini 2000; Kauffmann Haehnelt 2000; Yu Tremaine 2002; Steed Weinberg 2003; Wyithe Loeb 2003; Granato et al. |
2004. 2006: Marconi et al. | 2004, 2006; Marconi et al. |
2004: Moerloni al. | 2004; Merloni et al. |
2004: Yu Lu 2004: Miralca-IEscucde Ixollmeier 2005: Murray et al. | 2004; Yu Lu 2004; Miralda-Escudè Kollmeier 2005; Murray et al. |
2005: Cattaneo et al. | 2005; Cattaneo et al. |
2006: Croton et al. | 2006; Croton et al. |
2006: Hopkins et al. | 2006; Hopkins et al. |
2006: Lapi et al. | 2006; Lapi et al. |
2006: Shankar et al. | 2006; Shankar et al. |
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