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. Placing these detections iu the context of the burst Traction among equally massive cvarf galaxies is difficult. i0Wweveor. since we cannot currently constrain their nunber density: galaxies without starbursts have older uumositv weighted ages and. as a cousequence. are at cast 2 maeuituces fainter than the starbursting ealaxics.
Placing these detections in the context of the burst fraction among equally massive dwarf galaxies is difficult, however, since we cannot currently constrain their number density; galaxies without starbursts have older luminosity weighted ages and, as a consequence, are at least 2 magnitudes fainter than the starbursting galaxies.
With marginal detections. even in the latest WECS3 cata. and no obvious spectral features. their redshifts caunot casily be estimated.
With marginal detections, even in the latest WFC3 data, and no obvious spectral features, their redshifts cannot easily be estimated.
Model predictions differ stronely roni observational measurements of the galaxw stellar nass function (Ciootal.2011). and caunot be used ο constrain the burst fraction among the population of ON-nuass galaxies.
Model predictions differ strongly from observational measurements of the galaxy stellar mass function \citep{guo11} and cannot be used to constrain the burst fraction among the population of low-mass galaxies.
Nevertheless. we can still gauge the importance of the observed burst for the formation of low-mass galaxies by uakiug the reasonable (aud testable) assuiuptiou that he observed bursts occur with the same frequency at all epochs 1So.835 a period of —1 Cyr durius which the cosimic star formation history peaked ancl after which the απο density of starbursts declines as nentioned above.
Nevertheless, we can still gauge the importance of the observed burst for the formation of low-mass galaxies by making the reasonable (and testable) assumption that the observed bursts occur with the same frequency at all epochs $1\lesssim z \lesssim 3.5$ – a period of $\sim$ 4 Gyr during which the cosmic star formation history peaked and after which the number density of starbursts declines as mentioned above.
The basic iudication that such bursts are Hmuportant is that the nuniber of stars produced im such bursts over a period of several Cor is comparable o the number stars in preseut-dav dwarf galaxies.
The basic indication that such bursts are important is that the number of stars produced in such bursts over a period of several Gyr is comparable to the number stars in present-day dwarf galaxies.
Let us use this consideration to construct a simple tov model hat relates the observations preseuted here to the mass Unetiou of present-day low-ass galaxies.
Let us use this consideration to construct a simple toy model that relates the observations presented here to the mass function of present-day low-mass galaxies.
Cowctal.(2011) use the data from Baldivctal.(2008). for the ealaxies with masses down to 101 AY... and their mass uction eau be represeuted by a simple power law forall galaxies with masses <1019 AL... that is. well below the shee of the Schechter(1976). function: in units of Mpc3dlog(M) 5.
\citet{guo11} use the data from \citet{baldry08} for the galaxies with masses down to $10^7~\msol$ , and their mass function can be represented by a simple power law forall galaxies with masses $<10^{10}~\msol$ , that is, well below the knee of the \citet{schechter76} function: in units of $\rm{Mpc}^{-3}~\rm{d}\log(M_*)^{-1}$ .
Let us then express the stellar mass of the preseut- descendants of the observed starburstiug galaxies at Dod] interms of the following star formation history:
Let us then express the stellar mass of the present-day descendants of the observed starbursting galaxies at $z\sim 1.7$ interms of the following star formation history:
Since the number of stars in any post-MS stage scales linearly with the total luminosity of the stellar population 1988).. this ratio is predicted to be equal to | for any (not segregated) reference populations. while a larger (smaller) value is expected for populations which are more (less) concentrated.
Since the number of stars in any post-MS stage scales linearly with the total luminosity of the stellar population , this ratio is predicted to be equal to 1 for any (not segregated) reference populations, while a larger (smaller) value is expected for populations which are more (less) concentrated.
In order to estimate the sampled light. we have integrated the King profile best-fitting the surface density distribution (from S11).
In order to estimate the sampled light, we have integrated the King profile best-fitting the surface density distribution (from S11).
The radial trend of μον and Rpss in the same annuli previously defined is shown in the lower panel of Fig. 4..
The radial trend of $R_{\rm RGB}$ and $R_{\rm BSS}$ in the same annuli previously defined is shown in the lower panel of Fig. \ref{radist}.
As apparent. it turns out to be constant and equal to | not only for the RGB (and the HB) stars. but also for the BSS. thus further demonstrating that this population has a radial distribution fully consistent with that of the reference ones.
As apparent, it turns out to be constant and equal to 1 not only for the RGB (and the HB) stars, but also for the BSS, thus further demonstrating that this population has a radial distribution fully consistent with that of the reference ones.
The radial distribution of BSS in PalId is indistinguishable from that of its normal (and less massive) stars.
The radial distribution of BSS in Pal14 is indistinguishable from that of its normal (and less massive) stars.
As in the case of the only two other clusters showing the same featurerespectively)... this is an observational proof that Pall4 is dynamically young. still far from having established energy equipartition even in its innermost regions.
As in the case of the only two other clusters showing the same feature, this is an observational proof that Pal14 is dynamically young, still far from having established energy equipartition even in its innermost regions.
This is in agreement with the extremely long half-mass radius relaxation time (~20 Gyr) recently estimated by SII.
This is in agreement with the extremely long half-mass radius relaxation time $\sim 20$ Gyr) recently estimated by S11.
Moreover. our results suggest that the unusually flat mass function. measured by cannot be explained by energy equipartition developed during the cluster dynamical evolution. but should be primordial2011).
Moreover, our results suggest that the unusually flat mass function measured by cannot be explained by energy equipartition developed during the cluster dynamical evolution, but should be primordial.
. We note that since the mass range covered by that study is quite limited. further investigation is needed.
We note that since the mass range covered by that study is quite limited, further investigation is needed.
Finally. the negligible degree of relaxation of Pall4 suggests that the observed tidal tails should not be preferentially populated by low mass stars evaporated from the cluster.
Finally, the negligible degree of relaxation of Pal14 suggests that the observed tidal tails should not be preferentially populated by low mass stars evaporated from the cluster.
The flat BSS radial distribution also suggests that (as expected in such a low density environment) stellar collisions played a minor role in generating these exotica and affecting the binary population.
The flat BSS radial distribution also suggests that (as expected in such a low density environment) stellar collisions played a minor role in generating these exotica and affecting the binary population.
Hence. as in the case of «Centauri and NGC 2419. the BSS we are observing likely derive from the evolution of primordial binaries and can be used to get a rough estimate of the fraction of such a population.
Hence, as in the case of $\omega$ Centauri and NGC 2419, the BSS we are observing likely derive from the evolution of primordial binaries and can be used to get a rough estimate of the fraction of such a population.
As a first consideration. we note that the number of BSS normalized to the sampled luminosity is Sdpss72 InwCentauri and S4pss73 in NGC 2419.
As a first consideration, we note that the number of BSS normalized to the sampled luminosity is $S4_{\rm BSS}\simeq 2$ in $\omega$ Centauri and $S4_{\rm BSS}\simeq 3$ in NGC 2419.
The same ratio in Pall4 rises to 29 (1.8. a value 10 times larger than what found in the two clusters with similar BSS radial distribution).
The same ratio in Pal14 rises to 29 (i.e. a value 10 times larger than what found in the two clusters with similar BSS radial distribution).
However this value is not that surprising when compared to the field.
However this value is not that surprising when compared to the field.
In fact. as discussed by(2006).. the observed BSS specific frequency Ngss/Nup=0.1 in «Centauri turned out to be ~40 times lower that what observed in the field (Ngss/Nug= 4) by Preston Sneden (2000).
In fact, as discussed by, the observed BSS specific frequency $N_{\rm BSS}/N_{\rm HB}=0.1$ in $\omega$ Centauri turned out to be $\sim 40$ times lower that what observed in the field $N_{\rm BSS}/N_{\rm HB}=4$ ) by Preston Sneden (2000).
The value found in Palld is Npss/Nap~1l. in much better agreement with the above field sample.
The value found in Pal14 is $N_{\rm BSS}/N_{\rm HB}\sim 1$, in much better agreement with the above field sample.
Under the hypothesis that all the BSS are originated by primordial binaries. this possibly suggests that the binary fraction in Pall4 (and in the field) might be much higher than in the other two GCs
Under the hypothesis that all the BSS are originated by primordial binaries, this possibly suggests that the binary fraction in Pal14 (and in the field) might be much higher than in the other two GCs.
Aroughestimateo fthebinary fractioninPalomarldcanbederivec
A rough estimate of the binary fraction in Palomar 14 can be derived from the correlation between the measured binary fraction and the cluster integrated magnitude found by in a sample of 18 open and low-density globular clusters.
(STI) we predict η~30-40% for Palomar 14.
From their Figure 7, and adopting $M_V=-4.95$ (S11) we predict $f_{\rm bin}\sim 30-40\%$ for Palomar 14.
In addition. the comparison between the number of BSS per unit luminosity and the fraction of binaries measured in a sample of low-density GCs (in. which the collisional channel of BSS formation is expected to be negligible) shows that while BSS- GCs (with 8<Sdpss 13) host a small fraction of
In addition, the comparison between the number of BSS per unit luminosity and the fraction of binaries measured in a sample of low-density GCs (in which the collisional channel of BSS formation is expected to be negligible) shows that while BSS-poor GCs (with $8<S4_{\rm BSS}<13$ ) host a small fraction of
NCS operation.
NCS operation.
However. none of the galaxies in the UDF reached signal levels requiring correction therefore linearity is not an issue in this analvsis.
However, none of the galaxies in the UDF reached signal levels requiring correction therefore linearity is not an issue in this analysis.
A cosmic ταν event produces a sharp jump in signal intensity in the first readout alter the event.
A cosmic ray event produces a sharp jump in signal intensity in the first readout after the event.
Most events do not saturate (he pixel aud subsequent readouts continue to monitor ihe incident flix.
Most events do not saturate the pixel and subsequent readouts continue to monitor the incident flux.
At this point in the analvsis the readouts are stored as delta signal levels between each readout.
At this point in the analysis the readouts are stored as delta signal levels between each readout.
The signal “ramp” is reconstructed by adding the deltas together.
The signal “ramp” is reconstructed by adding the deltas together.
The first step in cosmic rav detection is a linear fit to the signal ramp. which will be a poor fit to the data if (here is a cosmic rav jump.
The first step in cosmic ray detection is a linear fit to the signal ramp, which will be a poor fit to the data if there is a cosmic ray jump.
The residuals to the fit will be increasingly. negative wilh a sharp transition {ο positive after the event.
The residuals to the fit will be increasingly negative with a sharp transition to positive after the event.
The cosmic ray detection procedure looks for the negative to positive transilion as a signature of a cosmic ταν hit.
The cosmic ray detection procedure looks for the negative to positive transition as a signature of a cosmic ray hit.
If it detects a residual transition above the level expected [rom noise it removes the delta signal between the (vo readouts before and after (he event. recaleulates (he signal ramp and fits a new linear solution.
If it detects a residual transition above the level expected from noise it removes the delta signal between the two readouts before and after the event, recalculates the signal ramp and fits a new linear solution.
The cosmic rav. procedure rechecks (he ramp to see if (here was another cosmic ray hit and removes (he proper delta signal if one is detected.
The cosmic ray procedure rechecks the ramp to see if there was another cosmic ray hit and removes the proper delta signal if one is detected.
If there is still a detectable cosmic ray signature after the second velit (he pixel is marked as bad aud no further correction attenmpl is made.
If there is still a detectable cosmic ray signature after the second refit the pixel is marked as bad and no further correction attempt is made.
If the signal is saturated after the cosmic rav hit only the signal obtained before the event is used in the analvsis.
If the signal is saturated after the cosmic ray hit only the signal obtained before the event is used in the analysis.
The final recorded signal is the value of the slope of the linear fit to the signal ramp in acus per second.
The final recorded signal is the value of the slope of the linear fit to the signal ramp in adus per second.
Cosmic ray hits (hat occur in the 0.3 seconds between (the first ancl second read are detected as fits that do not intercept zero. recalling that the first read is subtracted [rom second so that the second read is the first point in the ramp.
Cosmic ray hits that occur in the 0.3 seconds between the first and second read are detected as fits that do not intercept zero, recalling that the first read is subtracted from second so that the second read is the first point in the ramp.
These hits do not affect the calculated slope but are marked as cosmic ray hits in the data quality array. discussed later.
These hits do not affect the calculated slope but are marked as cosmic ray hits in the data quality array discussed later.
Each quadrant of the NICMOS detectors has a separate output amplifier to transmit ihe analog signal to an A/D converter.
Each quadrant of the NICMOS detectors has a separate output amplifier to transmit the analog signal to an A/D converter.
This was done to prevent the loss of an entire detector θαΤαν if (here was a [failure of an output amplifier.
This was done to prevent the loss of an entire detector array if there was a failure of an output amplifier.
As a side elfect of this design decision a small DC bias offset can occur between the 4 detector quadrants.
As a side effect of this design decision a small DC bias offset can occur between the 4 detector quadrants.
Although the olfset is small. the effect is significant relative to the faint galaxy signals in the UDF and can cause unreliable outputs during the drizzle process.
Although the offset is small, the effect is significant relative to the faint galaxy signals in the UDF and can cause unreliable outputs during the drizzle process.
Since (here is significant skv background from
Since there is significant sky background from
distribution of thehadron £i
in $p_T^2$ space where $m_H$ denotes the mass of $H$ .
nthe cascade processes isgiven bythe distribution function vo». ν πμ a 3.
The parameter $\alpha$ is fixed to $\alpha=1.8$ $^{-\frac{3}{2}}$ by using experimental data on ${\mbox{\boldmath$ $}_{T}^2}$ distributionsof pions in $\pi p$ Details
old. we detect no dilference in their mass- and. luminosity-weighted ages.
old, we detect no difference in their mass- and luminosity-weighted ages.
On the contrary. luminositv-weightec ages [or discs are 2.S] Gyr. significantly smaller than the mass-weighted estimates.
On the contrary, luminosity-weighted ages for discs are $2-8$ ] Gyr, significantly smaller than the mass-weighted estimates.
The mean ages of our simulated. disces and spheroids are in good agreement with observational results.
The mean ages of our simulated discs and spheroids are in good agreement with observational results.
We find. clear. signs of the presence of more than one component in our simulated clises. reminiscent of observed thin and thick cises.
We find clear signs of the presence of more than one component in our simulated discs, reminiscent of observed thin and thick discs.
The voungest stars define thin structures. with high tangential velocities and low velocity dispersions. whereas the oldest. dise stars celine thicker discs. with lower rotational velocities and uigher velocity dispersions.
The youngest stars define thin structures, with high tangential velocities and low velocity dispersions, whereas the oldest disc stars define thicker discs, with lower rotational velocities and higher velocity dispersions.
Assuming that thin ancl thick disc components can be distinguished by the age of their stars (we adopted 9 Cyr as the boundary). we determined vpical rotation velocities. and velocity dispersions for hick/thin simulated discs.
Assuming that thin and thick disc components can be distinguished by the age of their stars (we adopted $9$ Gyr as the boundary), we determined typical rotation velocities and velocity dispersions for thick/thin simulated discs.
Phese agree reasonably well with observations of the Milkv Way.
These agree reasonably well with observations of the Milky Way.
We note. however. that our simulations cover a range of galaxy masses and halo ooperties. and. therefore of tvpical rotation velocities and velocity clispersions.
We note, however, that our simulations cover a range of galaxy masses and halo properties, and therefore of typical rotation velocities and velocity dispersions.
As expected in the context of the AC'DAL cosmology assumed in. this work. the stellar components of our eight simulations show great 7variety dn their structure.
As expected in the context of the $\Lambda$ CDM cosmology assumed in this work, the stellar components of our eight simulations show great variety in their structure.
Discs have a wide range of —licknesses and scale-Ieneths. and are usually complex. with misaligned components or boxy shape.
Discs have a wide range of thicknesses and scale-lengths, and are usually complex, with misaligned components or boxy shape.
The inner regions of spheroids also show great diversity: half of the simulated galaxies have bar components. which can dominate over the bulges.
The inner regions of spheroids also show great diversity: half of the simulated galaxies have bar components, which can dominate over the bulges.
Sérrsic fits to the mass density profiles generally give shape parameters n~1L. i.e. they are similar to exponential.
Sérrsic fits to the mass density profiles generally give shape parameters $n\sim 1$, i.e. they are similar to exponential.
As found in other studies. outer spheroids or "stellar haloes” are very rich in structure. with streams. clumps and shells. similar to observational results (Alartinnez-Deleaco et al.
As found in other studies, outer spheroids or “stellar haloes” are very rich in structure, with streams, clumps and shells, similar to observational results nez-Delgado et al.
2010).
2010).
In broad terms. the sizes and structure of our bulges and disces agree well with observational results. although our bulges are too massive with respect to discs.
In broad terms, the sizes and structure of our bulges and discs agree well with observational results, although our bulges are too massive with respect to discs.
Vhese are useful to investigate the domünant formation channels for the cdillerent stellar components.
These are useful to investigate the dominant formation channels for the different stellar components.
We find that discs have the highestdn-5u fractions. typically z0.9. and all dise stars vounger than 9 Gyr formed in the dise itself.
We find that discs have the highest fractions, typically $\gtrsim 0.9$, and all disc stars younger than $9$ Gyr formed in the disc itself.
la one of the simulated ealaxies. however. 15% of the final disc mass is contributed by à satellite that came in on a nearly coplanar orbit.
In one of the simulated galaxies, however, $15\%$ of the final disc mass is contributed by a satellite that came in on a nearly coplanar orbit.
These stars are generally old but are able to stay in the disc.
These stars are generally old but are able to stay in the disc.
Inner spheroids (1.0. bulges and bars) also have relatively high fractions. generally larger than 0.8. and consequently a low contribution from accreted stars formed. in systems other than the main progenitor.
Inner spheroids (i.e. bulges and bars) also have relatively high fractions, generally larger than $0.8$, and consequently a low contribution from accreted stars formed in systems other than the main progenitor.
Conversely. outer spheroids have a large contribution [rom acerctec stars ancl low fractions. in the range 0.151.35.
Conversely, outer spheroids have a large contribution from accreted stars and low fractions, in the range $0.15-0.35$.
As expected. the lowest. values are detected. for svstems with recent massive mergers.
As expected, the lowest values are detected for systems with recent massive mergers.
Our results suggest that the outer regions of galaxies are populated. by a combination ofzn-sihi and accreted stars. with the relative fraction. rellecting the particular formation and accretion history of the host.
Our results suggest that the outer regions of galaxies are populated by a combination of and accreted stars, with the relative fraction reflecting the particular formation and accretion history of the host.
We also find a negative dependence of fraction on radius. the outermost regions of simulated galaxies are almost entirely populated: by accreted stars.
We also find a negative dependence of fraction on radius, the outermost regions of simulated galaxies are almost entirely populated by accreted stars.
These results are in &ood agreement with other studies (e.g. Zolotov et al.
These results are in good agreement with other studies (e.g. Zolotov et al.
2009: Cooper et al.
2009; Cooper et al.
2010).
2010).
We [ind complex and. highly asvmmetric dvnamical structures for the stellar and gaseous components of our simulated galaxies.
We find complex and highly asymmetric dynamical structures for the stellar and gaseous components of our simulated galaxies.
They rellect not only the diversity of our discs. bulges and bars. but also the non-trivial interplay between cooling and heating mechanisms.
They reflect not only the diversity of our discs, bulges and bars, but also the non-trivial interplay between cooling and heating mechanisms.
Circular velocity curves are relatively Hat for à5 κρὸ and have peak velocities which are 10 to 20% higher than the value at [arger raclii.
Circular velocity curves are relatively flat for $r\gtrsim 5$ kpc, and have peak velocities which are $10$ to $20\%$ higher than the value at larger radii.
We tested the elfects of resolution by comparing the results of our simulations to several additional simulations of 4 ancl S times lower mass resolution.
We tested the effects of resolution by comparing the results of our simulations to several additional simulations of $4$ and $8$ times lower mass resolution.
We find relatively good agreement for the formation time-scales. structure ancl fractions of cliscs and spheroics.
We find relatively good agreement for the formation time-scales, structure and fractions of discs and spheroids.
The ἄνπαίσα properties show more significant variation with resolution. in particular for the discs.
The dynamical properties show more significant variation with resolution, in particular for the discs.
We find that in simulations of higher resolution the disces tend to have higher tangential velocities and lower velocity dispersions.
We find that in simulations of higher resolution the discs tend to have higher tangential velocities and lower velocity dispersions.
These. results indicate that low resolution runs suller from. artificial clise heating.
These results indicate that low resolution runs suffer from artificial disc heating.
However. the trends found for the velocities and velocity. dispersions with age are still captured in the Low resolution runs.
However, the trends found for the velocities and velocity dispersions with age are still captured in the low resolution runs.
In broad terms. our simulations agree with previous studies. and. also show similar deficiencies. the most important one being the presence of overly dominant bulges.
In broad terms, our simulations agree with previous studies, and also show similar deficiencies, the most important one being the presence of overly dominant bulges.
Lt is still under debate whether these deficiencies are due to insullicient resolution. or are à consequence of a poor description of the eas hydrodynamies or. the star formation and feedback processes.
It is still under debate whether these deficiencies are due to insufficient resolution, or are a consequence of a poor description of the gas hydrodynamics or the star formation and feedback processes.
On the other hand. the persistent [ailure to reproduce late-twpe galaxies in a cosmological context in typical galactic haloes may indicate that additional physical processes. not vet considered. in simulations. play a role in the regulation of star formation in ealaxies.
On the other hand, the persistent failure to reproduce late-type galaxies in a cosmological context in typical galactic haloes may indicate that additional physical processes, not yet considered in simulations, play a role in the regulation of star formation in galaxies.
Our understanding of important. processes related to galaxy formation — from the formation of stars to the elfects. of supernova and black hole. feedback is still quite rudimentary. so it is remarkable that cosmological simulations nevertheless produce galaxies with structural ancl dynamical properties in reasonable agreement with observational results.
Our understanding of important processes related to galaxy formation – from the formation of stars to the effects of supernova and black hole feedback – is still quite rudimentary, so it is remarkable that cosmological simulations nevertheless produce galaxies with structural and dynamical properties in reasonable agreement with observational results.
We thank the referee for a thorough reading of this work ancl for his/her helpful. comments ancl suggestions.
We thank the referee for a thorough reading of this work and for his/her helpful comments and suggestions.
The simulations were carried out at the Computing Centre of the Max-Planck-Society in Garching.
The simulations were carried out at the Computing Centre of the Max-Planck-Society in Garching.
This research was supported bv the DEC cluster of excellence. "Origin. and Structure of the Universe’.
This research was supported by the DFG cluster of excellence `Origin and Structure of the Universe'.
CS thanks M. Bovlan-IExolchin for providing the data for Fig. Αι.
CS thanks M. Boylan-Kolchin for providing the data for Fig. \ref{MAH},
Lauren MacArthur for her help in the interpretation of observational results. and A. Steinmetz. M. Vlajic. J. Wang and M. Williams for useful and stimulating discussions.
Lauren MacArthur for her help in the interpretation of observational results, and M. Steinmetz, M. Vlajic, J. Wang and M. Williams for useful and stimulating discussions.
in many crucial aspects from that of the other classes of ACN; they lack the substantial BLR. tori and thermal dise enission. which are usually associated with active nuclei. and sugeest that in FR Is vaccretion might take place in a low efficiency. radiative regime.
in many crucial aspects from that of the other classes of AGN; they lack the substantial BLR, tori and thermal disc emission, which are usually associated with active nuclei”, and suggest that in FR Is “accretion might take place in a low efficiency radiative regime”.
Areguneuts which try to explain the FR ΤΕΙ II dichotomy based solely on differences of the ceutral cneine are not free from difficulties: first. the similar properties between parsec scale jets in FR IT and FR ILradio galaxies inplv the same (or very similar) jet production mechanism for both types of sources. and second. the existence of FR I/II sources with large scale properties common with FR I and FR II radio galaxies (Copalsvishua Wiita 2001)).
Arguments which try to explain the FR I-FR II dichotomy based solely on differences of the central engine are not free from difficulties: first, the similar properties between parsec scale jets in FR I and FR II radio galaxies imply the same (or very similar) jet production mechanism for both types of sources, and second, the existence of FR I/II sources with large scale properties common with FR I and FR II radio galaxies (Gopal-Krishna Wiita \cite{gopal2}) ).
We support the idea that the central kiloparsec reeion surrounding the active core ust also play. together with the central cugine. a crucial role iu the FR LFR IL dichotomy.
We support the idea that the central kiloparsec region surrounding the active core must also play, together with the central engine, a crucial role in the FR I-FR II dichotomy.