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We now turn our attention to the spatial distribution of the satellite galaxies of the MW.
We now turn our attention to the spatial distribution of the satellite galaxies of the MW.
Figure 2 shows the corresponding y-¢ plot for the 11 most luminous satellites of the Milky Way iidentical to Metzetal. (2007))).
Figure \ref{figure:gammazeta} shows the corresponding $\gamma$ $\zeta$ plot for the 11 most luminous satellites of the Milky Way identical to \citet{Metz07}) ).
The significance, orientation and thickness of a plane to the classical dwarf galaxy sample is given in Table 1..
The significance, orientation and thickness of a plane to the classical dwarf galaxy sample is given in Table \ref{table:lsq}.
Our results may be compared to those of Metzetal.(2007) who find @=158:2, b=—11:9, and a rms thickness, A=32.6 kpc.
Our results may be compared to those of \citet{Metz07} who find $\ell =158\mathring{.}2$, $\textit{b} =-11\mathring{.}9$, and a rms thickness, $\Delta$ =32.6 kpc.
Figure 2 (center) shows our findings with the inclusion of 10 recently discovered MW satellites (replicating the sample of Metzetal. (2009b))).
Figure \ref{figure:gammazeta} (center) shows our findings with the inclusion of 10 recently discovered low-luminosity MW satellites (replicating the sample of \citet{Metz09}) ).
Here our results may be compared to those of Metzetal.(2009b) who find /=159°74+2°3, b=—6°8+2°3, and A=24.9+1.1 kpc.
Here our results may be compared to those of \citet{Metz09} who find $\ell = 159\mathring{.}7 \pm 2\mathring{.}3$, $\textit{b} = -6\mathring{.}8 \pm 2\mathring{.}3$, and $\Delta$ $\pm 1.1$ kpc.
It may be argued that on the basis of the results of Beslaetal.(2010) that the Magellanic Clouds (MCs) be considered separately since their proper motions suggest they are on their first passage by the MW.
It may be argued that on the basis of the results of \citet{Besla10} that the Magellanic Clouds (MCs) be considered separately since their proper motions suggest they are on their first passage by the MW.
On the other hand, since the MCs lie on a plane common to the remaining 9 'classical satellites we propose that if the MCs are on their first infall that they are further evidence in favor of preferred accretion along a common plane.
On the other hand, since the MCs lie on a plane common to the remaining 9 'classical' satellites we propose that if the MCs are on their first infall that they are further evidence in favor of preferred accretion along a common plane.
In both of the above samples we recover the significance, orientation and thickness of the planar distributions discussed in the literature.
In both of the above samples we recover the significance, orientation and thickness of the planar distributions discussed in the literature.
It is apparent from Table 1 that the orientation and thickness of the plane that describes the spatial distribution of the YH GCs at Roo>10 kpc is indistinguishable from that of the classical MW dwarf galaxies and the recently discovered ‘ultra-faint’ dwarf galaxies.
It is apparent from Table \ref{table:lsq} that the orientation and thickness of the plane that describes the spatial distribution of the YH GCs at $R_{GC} > 10$ kpc is indistinguishable from that of the classical MW dwarf galaxies and the recently discovered `ultra-faint' dwarf galaxies.
For this reason we combine these three samples to define a common plane of satellites (PoS).
For this reason we combine these three samples to define a common plane of satellites (PoS).
Figure 3 shows the spatial distribution on the sky of the classes of object we have considered.
Figure \ref{figure:PoS} shows the spatial distribution on the sky of the classes of object we have considered.
The PoS, as defined above, is shown as the solid line.
The PoS, as defined above, is shown as the solid line.
The grayscale probability density function seen at /~156° shows the positions of the normal to this plane in 10* random realisations of the input catalogue distance uncertainties.
The grayscale probability density function seen at $\ell \sim 156^{\circ}$ shows the positions of the normal to this plane in $10^4$ random realisations of the input catalogue distance uncertainties.
The top left panel of Figure 4 shows a section through the vicinity of the MW as seen at a viewing angle edge on to our derived PoS. We note that the OH GC NGC 2419 (the open
The top left panel of Figure \ref{figure:PoS_slices} shows a section through the vicinity of the MW as seen at a viewing angle edge on to our derived PoS. We note that the OH GC NGC 2419 (the open
quantitative description of the properties of elusive high-redshift, possibly primordial, galaxies.
quantitative description of the properties of elusive high-redshift, possibly primordial, galaxies.
They are also very useful to interpret the data coming from deep surveys as the HST/WFC3 and future ones.
They are also very useful to interpret the data coming from deep surveys as the HST/WFC3 and future ones.
However, there is considerable room for improvement left by our study.
However, there is considerable room for improvement left by our study.
In the following we would like to elaborate on the uncertainties and shortcomings of our findings.
In the following we would like to elaborate on the uncertainties and shortcomings of our findings.
We first note that resolving the dwarf galaxy population and following the PopIII transition process along with the large variety of physical processes implemented in the simulation limits the size of the cosmic volume that can be simulated.
We first note that resolving the dwarf galaxy population and following the PopIII transition process along with the large variety of physical processes implemented in the simulation limits the size of the cosmic volume that can be simulated.
Resolution is certainly an important issue, as it is well known (see, e.g. Governato et al.
Resolution is certainly an important issue, as it is well known (see, e.g. Governato et al.
2010) to affect the simulated star formation rates and cause the loss of sub-galactic structures.
2010) to affect the simulated star formation rates and cause the loss of sub-galactic structures.
The dependence of the results from resolution has been presented an analyzed for the same set-up of the present simulations in Tornatore et al. (
The dependence of the results from resolution has been presented an analyzed for the same set-up of the present simulations in Tornatore et al. (
2007) to which we refer the interested reader (see Fig.
2007) to which we refer the interested reader (see Fig.
1 of that paper).
1 of that paper).
Resolution might also alter the details of the "Pop III wave" evolution, since galactic substructure allows star formation to occur at the edges of galaxies as well.
Resolution might also alter the details of the "Pop III wave" evolution, since galactic substructure allows star formation to occur at the edges of galaxies as well.
The PopIII-PopII transition is also very dependent on the assumed IMF of PoplIII stars.
The PopIII-PopII transition is also very dependent on the assumed IMF of PopIII stars.
Our conclusions are valid under the assumption that PopIII stars were very massive (M>100 Mc) and the first metal production is driven by the explosion of pair-instability SNe (see Schneider et al.
Our conclusions are valid under the assumption that PopIII stars were very massive $M\ge 100\;\Msun$ ) and the first metal production is driven by the explosion of pair-instability SNe (see Schneider et al.
2006 for alternatives).
2006 for alternatives).
Our box is also too small to properly describe cosmic reionization, let alone that we are not even attempting to properly treat radiative transfer.
Our box is also too small to properly describe cosmic reionization, let alone that we are not even attempting to properly treat radiative transfer.
These issues have been already addressed in previous works of our group; as already stated we are concerned here with the properties of high-z galaxies, which are presumably more affected by their internal physics rather than by the environment, as we explain below.
These issues have been already addressed in previous works of our group; as already stated we are concerned here with the properties of high-z galaxies, which are presumably more affected by their internal physics rather than by the environment, as we explain below.
Comparing the simulated volumes with the observed ones is very challenging as considerable uncertainty exists on the latter (see discussion in Appendix B of Bouwens et al.
Comparing the simulated volumes with the observed ones is very challenging as considerable uncertainty exists on the latter (see discussion in Appendix B of Bouwens et al.
2009).
2009).
However, as it could be induced from the extension of the LF towards the most luminous and rarest objects at z=7, we estimate that the volume sampled by experiments should be about 30 times larger than our simulated one.
However, as it could be induced from the extension of the LF towards the most luminous and rarest objects at z=7, we estimate that the volume sampled by experiments should be about 30 times larger than our simulated one.
The next caveat comes from the fact that we have neglected the effects of minihalos (virial temperature «10* K).
The next caveat comes from the fact that we have neglected the effects of minihalos (virial temperature $< 10^4$ K).
Our resolution does not allow us to track the formation of such objects, whose stellar contribution remains very uncertain due to radiative feedback effects (Haiman Bryan 2006; Susa Umemura 2006; Ahn Shapiro 2007, Okamoto, Gao Theuns 2008, Salvadori Ferrara 2009).
Our resolution does not allow us to track the formation of such objects, whose stellar contribution remains very uncertain due to radiative feedback effects (Haiman Bryan 2006; Susa Umemura 2006; Ahn Shapiro 2007, Okamoto, Gao Theuns 2008, Salvadori Ferrara 2009).
The presence of such small collapsed structures, if able to form stars, could alter the reionization history to some extent and increase the number counts of high-redshift galaxies, if detectable.
The presence of such small collapsed structures, if able to form stars, could alter the reionization history to some extent and increase the number counts of high-redshift galaxies, if detectable.
As far as reionization is concerned, it has already been shown by Choudhury Ferrara (2007) that acceptable fits to all relevant reionization data can be obtained without any need for PopllI stars.
As far as reionization is concerned, it has already been shown by Choudhury Ferrara (2007) that acceptable fits to all relevant reionization data can be obtained without any need for PopIII stars.
The bulk of the ionizing photons in those models is produced by halos with virial temperatures just above 104 K, with increasingly better solution if normal, PoplI stars are allowed to form in minihalos.
The bulk of the ionizing photons in those models is produced by halos with virial temperatures just above $10^4$ K, with increasingly better solution if normal, PopII stars are allowed to form in minihalos.
As shown in Fig. 10,,
As shown in Fig. \ref{fig:mhalo},
where we plot the UV magnitude as function of the total mass of z=7 galaxies, sources detectable with JWST have halo masses10°Mo.
where we plot the UV magnitude as function of the total mass of z=7 galaxies, sources detectable with JWST have halo masses.
This corresponds to circular velocities υ.250 km s!.
This corresponds to circular velocities $v_c \simgt 50$ km $^{-1}$.
These objects are large enough that suppression by UVB photoionization filtering is at best marginal, if not negligible at all, as most of the works above agree upon.
These objects are large enough that suppression by UVB photoionization filtering is at best marginal, if not negligible at all, as most of the works above agree upon.
Internal mechanical feedback might be indeed more important.
Internal mechanical feedback might be indeed more important.
This process is however is already included at best in the simulations when computing the star formation rate of individual galaxies, modulo the many uncertainties that still plague our understanding of such phenomenon.
This process is however is already included at best in the simulations when computing the star formation rate of individual galaxies, modulo the many uncertainties that still plague our understanding of such phenomenon.
By using high resolution simulations specifically crafted to include the relevant physics of galaxy formation, along with a novel treatment of the metal dispersion that allows us to follow the PopIII-PoplI transition as dictated by the critical metallicity scenario, we have been able to reproduce the observed UV ΤΕΕ over a wide redshift range, 5<z«10.
By using high resolution simulations specifically crafted to include the relevant physics of galaxy formation, along with a novel treatment of the metal dispersion that allows us to follow the PopIII-PopII transition as dictated by the critical metallicity scenario, we have been able to reproduce the observed UV LFs over a wide redshift range, $5 < z < 10$.
We have also shown, by combining the simulation outputs with a dust model previously developed for LAEs (Dayal et al.
We have also shown, by combining the simulation outputs with a dust model previously developed for LAEs (Dayal et al.
2010), that dust effects at z©7 should be marginal, although this statement depends on many details (the dust properties and the dust distribution scale in the ISM) that are poorly constrained at this time.
2010), that dust effects at $z\approx 7$ should be marginal, although this statement depends on many details (the dust properties and the dust distribution scale in the ISM) that are poorly constrained at this time.
The general picture that can be drawn from our investigation is broadly consistent with the available data and therefore can be used to make specific predictions for the JWST.
The general picture that can be drawn from our investigation is broadly consistent with the available data and therefore can be used to make specific predictions for the JWST.
It is then useful to schematically summarize the main findings of the present work:
It is then useful to schematically summarize the main findings of the present work:
white dwarf not clissimilar to the U Sco type of recurrent novae. and that a search in plate archives for missed previous outbursts could. pay clivictoncls.
white dwarf, not dissimilar to the U Sco type of recurrent novae, and that a search in plate archives for missed previous outbursts could pay dividends.
A report on an carly N-rav. detection ancl following evolution of V2672 Oph was provided by Schwarzetal. (20090).
A report on an early X-ray detection and following evolution of V2672 Oph was provided by \citet{SOP09}.
.. Their. August 17.948 observation with theSwi satellite detected the nova with both the XIXE and UVOT instruments.
Their August 17.948 observation with the satellite detected the nova with both the XRT and UVOT instruments.
Further observations on the following. days found the nova to emit at a stable X-ray flux level. while rapicly declining at ultraviolet wavelengths.
Further observations on the following days found the nova to emit at a stable X-ray flux level, while rapidly declining at ultraviolet wavelengths.
Schwarzctal.(2000) suggested that the carly hare N-ray. emission. was likely due to shocks between the fast ejecta ancl a pre-existing circumstellar medium (as in the recurrent nova LS Oph e.g. Bodectal. 2006)) or intra-cjecta shocks (as3ode 1994).. Schwar
\citet{SOP09} suggested that the early hard X-ray emission was likely due to shocks between the fast ejecta and a pre-existing circumstellar medium (as in the recurrent nova RS Oph $-$ e.g. \citealt{BOO06}) ) or intra-ejecta shocks \citep[as in the very fast classical nova V838 Her,][]{OLB94}.
zetal.(2009) also looked for past X-ray observations in the field. and concluded that V2672 Oph was not recorded as ai X-ray source prior to the 2009 outburst.
\citet{SOP09} also looked for past X-ray observations in the field, and concluded that V2672 Oph was not recorded as an X-ray source prior to the 2009 outburst.
They also noted how the nova was not detected in gamma-ravs by IN'TTEGIUCAL/LIDBIS during Galactic bulge monitoring observations taken on 2009 August 20 and 23/24.
They also noted how the nova was not detected in gamma-rays by INTEGRAL/IBIS during Galactic bulge monitoring observations taken on 2009 August 20 and 23/24.
KraussHartman.Rupen.&Aliocluszewski(2009) reported their detection of radio cussion from V2672 Oph during the first two weeks after optical maximum.
\citet{HRM09} reported their detection of radio emission from V2672 Oph during the first two weeks after optical maximum.
The radio emission from most novae is dominated by thermal bremsstrahlung which is optically thick at early times. leading to a dependence of the Dux on frequeney as v “owherea~ 1 2 at carly times.
The radio emission from most novae is dominated by thermal bremsstrahlung which is optically thick at early times, leading to a dependence of the flux on frequency as $F_\nu \propto \nu^{+\alpha}$ , where $\alpha \sim$ 1 $-$ 2 at early times.
Vhev detected V2672 Oph at 8.46 Giz with the VLA on Sept 1.13. but obtained no detection at 2246 Cllz two days later. on September 3.18.
They detected V2672 Oph at 8.46 GHz with the VLA on Sept 1.13, but obtained no detection at 22.46 GHz two days later, on September 3.18.
IxraussHartman.Rupen.&Aliocuszewski(2009) concluded: that this was best explained. by a synchrotron origin for the radio emission observed from V2672 Oph.
\citet{HRM09} concluded that this was best explained by a synchrotron origin for the radio emission observed from V2672 Oph.
To support this view. they noted. that (7)) the strong shocks in its ejecta. suggested. by the hard. X-ray. emission. can also be the source for the relativistic electrons ancl strong magnetic fields needed to generate svnchrotron radiation. and that (11 1) the recurrent nova RS Oph. shows strong radio svnchrotron emission within clays of the outburst (e.g. 2009).
To support this view, they noted that $i$ ) the strong shocks in its ejecta, suggested by the hard X-ray emission, can also be the source for the relativistic electrons and strong magnetic fields needed to generate synchrotron radiation, and that $ii$ ) the recurrent nova RS Oph, shows strong radio synchrotron emission within days of the outburst \citep[e.g.,][]{PDB85,OBP06,EOB09}.
. The peculiarity and rarity of the phenomena displayed bv V2672 Oph is evident [rom these carb preliminary accounts.
The peculiarity and rarity of the phenomena displayed by V2672 Oph is evident from these early preliminary accounts.
In this paper we report our optical observations. that have allowed us. to. derive. the photometric and spectroscopic evolution and the basic properties of V2672 Oph. and which are used to perform. morpho-kincmatical modelling of the emission line profiles and hence disentangle the basic components of the expanding ejecta.
In this paper we report our optical observations, that have allowed us to derive the photometric and spectroscopic evolution and the basic properties of V2672 Oph, and which are used to perform morpho-kinematical modelling of the emission line profiles and hence disentangle the basic components of the expanding ejecta.
Photometric observations of V2672 Oph have been obtained with two instruments in collaboration with ο. Dallaporta. A. rigo. X. Siviero. S. Tomaselli. A. Maitan and 8S. Moretti of ANS Collaboration.
Photometric observations of V2672 Oph have been obtained with two instruments in collaboration with S. Dallaporta, A. Frigo, A. Siviero, S. Tomaselli, A. Maitan and S. Moretti of ANS Collaboration.
The first is a 0.25-m. Meade LX-200 Schmidt-Cassegrain. telescope located in Cembra (Trento. Παν).
The first is a 0.25-m Meade LX-200 Schmidt-Cassegrain telescope located in Cembra (Trento, Italy).
Lt is equipped with an SBIG ST-s CCD camera. 1020 array. 9 jam pixels — 0.74" /pix. with a Ποιά of view of 19 137.
It is equipped with an SBIG ST-8 CCD camera, $\times$ 1020 array, 9 $\mu$ m pixels $\equiv$ $^{\prime\prime}$ /pix, with a field of view of $^\prime$$\times$ $^\prime$.
The DVRede filters are [rom Schuler.
The $B$$V$$R_{\rm C}$$I_{\rm C}$ filters are from Schuler.
Vhe other one is a 0.25-m (/6 BHitchey-Chretien. robotic telescope. part of the GRAS network (GRASIS. Australia).
The other one is a 0.25-m f/6 Ritchey-Chretien robotic telescope, part of the GRAS network (GRAS15, Australia).
]t carries an ο). ST-LONATE CCD camera 1472 array. 6.S pum pixels — 0.93’ /pix. with a Ποια of view of 34.23)
It carries an SBIG ST-10XME CCD camera $\times$ 1472 array, 6.8 $\mu$ m pixels $\equiv$ $^{\prime\prime}$ /pix, with a field of view of $^\prime$$\times$ $^\prime$.
The BVRede filters are again from Schuler.
The $B$$V$$R_{\rm C}$$I_{\rm C}$ filters are again from Schuler.
The calibration of photometric zero points ancl colour equations have been carried out against the Landolt(1983.1992.2009) equatorial standards.
The calibration of photometric zero points and colour equations have been carried out against the \citet{L83,L92,L09} equatorial standards.
Phe photometric data are presented in Table 1.
The photometric data are presented in Table 1.
The total error. budget: (dominated. by the Poissonian noise. with only a minor contribution from the uncertainty of the transformation to the Landolt. system of cquatorial standards) does not exceed 0.035 mag for all points.
The total error budget (dominated by the Poissonian noise, with only a minor contribution from the uncertainty of the transformation to the Landolt system of equatorial standards) does not exceed 0.035 mag for all points.
Spectroscopic observations of V2672 Oph were obtained with two telescopes. in collaboration with D. Valisa. V. Luppi and P. Ochner of ANS Collaboration.
Spectroscopic observations of V2672 Oph were obtained with two telescopes, in collaboration with P. Valisa, V. Luppi and P. Ochner of ANS Collaboration.
A journal of the observations is given in Table 2.
A journal of the observations is given in Table 2.
The 0.6m. telescope of the Schiaparelli observatory in Varese (Italy). is equipped with a multimode spectrograph (Echelle | single dispersion modes) and a SBIG STIO-NAIL CCD camera 1472 array. 6.8 pe pixels).
The 0.6m telescope of the Schiaparelli observatory in Varese (Italy), is equipped with a multi-mode spectrograph (Echelle + single dispersion modes) and a SBIG ST10-XME CCD camera $\times$ 1472 array, 6.8 $\mu$ m pixels).
Ht was used to obtain low resolution. wide wavelength range spectra.
It was used to obtain low resolution, wide wavelength range spectra.
Ehe Asiago 1.22m telescope feces light to à DB&C€ spectrograph. equipped with a 1200 In/mm erating and XNDOItiDus 440A CCD camera (EEV 42-10BU back illuminated. chip. 512 pixels. 13.5 sam in size).
The Asiago 1.22m telescope feeds light to a C spectrograph, equipped with a 1200 ln/mm grating and ANDOR iDus 440A CCD camera (EEV 42-10BU back illuminated chip, $\times$ 512 pixels, 13.5 $\mu$ m in size).
Lt was used to obtain the higher resolution spectra around Ho.
It was used to obtain the higher resolution spectra around $\alpha$.
Phe spectra collected with both telescopes were calibrated. into absolute Duxes. by observations of several spectrophotometric standards. which were observed at air-masses close to those of V2672 Oph.
The spectra collected with both telescopes were calibrated into absolute fluxes by observations of several spectrophotometric standards, which were observed at air-masses close to those of V2672 Oph.
The zero points and slopes of the absolutely [uxed spectra were checked against the photometry of Table 1. by integrating the spectral Dux through the BVe photometric bands.
The zero points and slopes of the absolutely fluxed spectra were checked against the photometry of Table 1, by integrating the spectral flux through the $B$$V$$R_{\rm C}$ photometric bands.
The error on the fluxes of our spectra turned. out not to exceed over the wavelength range covered.
The error on the fluxes of our spectra turned out not to exceed over the wavelength range covered.
The photometric evolution of V2672 Oph is presented. in Figure 1.
The photometric evolution of V2672 Oph is presented in Figure 1.
The light- and colour-curves of Figure 1 were obtained by combining our data in Table 1 with other sources as follows.
The light- and colour-curves of Figure 1 were obtained by combining our data in Table 1 with other sources as follows.
The unfiltered CCD photometry obtained around. the time of discovery by various Japanese observers and reported bv Nakano.Yamaoka.&Ixadota (2009)... was scaled to V. band by adding |1.35 magnitudes.
The unfiltered CCD photometry obtained around the time of discovery by various Japanese observers and reported by \citet{nakano}, , was scaled to $V$ band by adding +1.35 magnitudes.
The Japanese amateurs
The Japanese amateurs
unity.
.
These offsets have been previously noted of Pleiades stars - see Fig.
These offsets have been previously noted of Pleiades stars - see Fig.
4 of "
4 of \cite{gorlova_spitzer_2006}.
Detection of an infrared excess is considered conclusive when the excess ratio of an object exceeds the upper 3c confidence interval of the entire distribution (?)).
Detection of an infrared excess is considered conclusive when the excess ratio of an object exceeds the upper $3\sigma$ confidence interval of the entire distribution \cite{sierchio10}) ).
Excluding AKII 437 and HII 132 (which significantly outlie), it is found that 7 objects (including 2 borderline objects) in the distribution satisfy this condition while falling to only 3 and 4 objects of the (?))
Excluding AKII 437 and HII 132 (which significantly outlie), it is found that 7 objects (including 2 borderline objects) in the distribution satisfy this condition while falling to only 3 and 4 objects of the \cite{sierchio10})
To estimate the mass loss from the cluster. we decided to use an ‘observational definition.
To estimate the mass loss from the cluster, we decided to use an `observational' definition.
At any A5given time we compare the cluster local density fgp, with the background5 stellar density fa. assuming that a star is actually belonging to the cluster i£ it is located in a region dense enough to make it distinguishable from the background. i.e. if being The limiting radius ry is then defined as the radius of the sphere (centered in the cluster density center) in whieh the cluster ‘emerges’ from the stellar background.
At any given time we compare the cluster local density $\rho_{gc}$ with the background stellar density $\rho_{\ast}$, assuming that a star is actually belonging to the cluster if it is located in a region dense enough to make it distinguishable from the background, i.e. if being The limiting radius $r_{L}$ is then defined as the radius of the sphere (centered in the cluster density center) in which the cluster `emerges' from the stellar background.
In Fig.25.. the evolution of the cluster mass. expressed in units of the initial mass Mj. is shown versus time for all the four simulations performed.
In \ref{mloss}, the evolution of the cluster mass, expressed in units of the initial mass $M_{0}$, is shown versus time for all the four simulations performed.
In the case of a cluster moving on orbils with apocenter <3.554. (he mass loss is dramatic: after about 30/4, the cluster loses about 75% of its mass: the best fit of the mass evolution as a function of time is given where / is expressed in units of (he bulge crossing (nme /,,,.
In the case of a cluster moving on orbits with apocenter $\leq 3.5 r_b$, the mass loss is dramatic: after about 30 $t_{\rm{cross}}$ the cluster loses about $\%$ of its mass; the best fit of the mass evolution as a function of time is given where $t$ is expressed in units of the bulge crossing time $t_{cross}$.
In (he remaining case. when the orbit extends up to 7.5 ry. the mass loss rate considerably diminishes and the cluster mass. after 30 ους. Is still about 6096 of its initial value.
In the remaining case, when the orbit extends up to 7.5 $r_b$, the mass loss rate considerably diminishes and the cluster mass, after 30 $t_{\rm{cross}}$ , is still about $\%$ of its initial value.
As is evident in Fig.25.. in this case (he mass loss rate increases every lime (he cluster passes at the minimum distance Iron the ealaxy center and not all particles which become unbound al perigalacticon are still so while moving again lo apogalacticon.
As is evident in \ref{mloss}, in this case the mass loss rate increases every time the cluster passes at the minimum distance from the galaxy center and not all particles which become unbound at perigalacticon are still so while moving again to apogalacticon.
It is possible to point out a region around the galaxy. center
It is possible to point out a region around the galaxy center
The present paper is the fourth of a series presenting the data accumulated by the public ESO Imaging Survey (EIS). being carried out in preparation for the first year of regular operation of the ΝΕΤ.
The present paper is the fourth of a series presenting the data accumulated by the public ESO Imaging Survey (EIS), being carried out in preparation for the first year of regular operation of the VLT.