ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
The observed-to-intrinsic luminosity ratio in the 2-10 keV band is then 0.675.! per cent for IRAS 0910444109.	The observed-to-intrinsic luminosity ratio in the 2–10 keV band is then $0.67\gamma^{-1}$ per cent for IRAS 09104+4109.
Tf the ratio of the 2-10 keV to bolometric luminosities is assumed to be {ηνyaar=0.05fixepaotoos-. typical for quasars (e.g. Elvis et al 1994). the bolometric luminosity is +.	If the ratio of the 2–10 keV to bolometric luminosities is assumed to be $f_{\rm HX/Bol}=0.05f_{\rm HX/Bol,0.05}$, typical for quasars (e.g., Elvis et al 1994), the bolometric luminosity is $1.5\times 10^{46}f_{\rm HX/Bol,0.05}^{-1} \gamma h^{-2}$ .
This is comparable to Li,=2.3 . and consistent. with. a picture. where the bulk of. the infrared emission is powered by a buried quasar through dust reradiation.	This is comparable to $L_{\rm ir}=2.3\times 10^{46}h^{-2}$ , and consistent with a picture where the bulk of the infrared emission is powered by a buried quasar through dust reradiation.
The radio source has an intermediate power between	The radio source has an intermediate power between
measurements of varying fidelity as a fuuetion of hour angle. a feature that must be incorporated in the combination of suapshots by a suitable welghtiug scheme.	measurements of varying fidelity as a function of hour angle, a feature that must be incorporated in the combination of snapshots by a suitable weighting scheme.
For each output image pixel we are attempting to form the best estimate of the sky polarized brightuess (by) from the combination of. η individual measurements (bjj.;). each of different weight (105:,:;).	For each output image pixel we are attempting to form the best estimate of the sky polarized brightness $\bar{\bf b}_{pix}$ ) from the combination of, $n$ individual measurements ${\bf b}_{pix,j}$ ), each of different weight $w_{pix,j}$ ).
The best such combined estimate is the weighted mean: where the weights (αρ) are generally the reciprocal of the pixel variance.	The best such combined estimate is the weighted mean: where the weights $w_{pix,j}$ ) are generally the reciprocal of the pixel variance.
The inverse-variauce of the MWA image pixels is best represented by the direction dependent polarimetric tile response iu the cirectiou of that pixel (the power respouse).	The inverse-variance of the MWA image pixels is best represented by the direction dependent polarimetric tile response in the direction of that pixel (the power response).
We are weightiug therelore by the power signal to noise ratio. uicer the assumption that the the thermal uoise is direction indepeudent.	We are weighting therefore by the power signal to noise ratio, under the assumption that the the thermal noise is direction independent.
Following the formalisin developed iu the series of papers: 2: 2: 2: ?. aud ο. the brightuess of au image pixel iu the measured. or instrument [rate is: where b are the coherency Lvectors. the superscript m iucicates “measured” aud should be taken to imply au instrument polarization basis (p. q) aud 5 represents a Stokes (4.Q.U. V) basis.	Following the formalism developed in the series of papers: \cite{hbs96}; \cite{shb96}; \cite{hb96}; \cite{ham00} and \cite{ham06}, the brightness of an image pixel in the measured, or instrument frame is: where ${\bf b}$ are the coherency 4-vectors, the superscript ${\bf m}$ indicates “measured” and should be taken to imply an instrument polarization basis $p,q$ ) and ${\bf s}$ represents a Stokes $I, Q, U, V$ ) basis.
The matrix 8 trausforms polarization from the instrument basis (orthogonal linear feeds) to the Stokes basis ou the sky aud is The MWA-32T is an interferometer comprising 32 dillerent elements but. as is described in §2.. each visibility is calibrated by au equalization process that allows us to approximate the iustruimenut polarimetric response by anmatrix.	The matrix $\stokes$ transforms polarization from the instrument basis (orthogonal linear feeds) to the Stokes basis on the sky and is The MWA-32T is an interferometer comprising 32 different elements but, as is described in \ref{sec:calibration}, each visibility is calibrated by an equalization process that allows us to approximate the instrument polarimetric response by an.
The 2x2 instrument Joues matrix J. determineW. how the generally non-orthiogonal projected receptors (p.q) respoud to the iucideut orthogoual sky polarization (X. Y).	The 2x2 instrument Jones matrix $\jones$, determines how the generally non-orthogonal projected receptors ) respond to the incident orthogonal sky polarization (X, Y).
As the projection of the tile receptors changes as a function of position the Jones matrix is different [or every pixel.	As the projection of the tile receptors changes as a function of position the Jones matrix is different for every pixel.
[tis this “instrument Joues matrix”. that is used to form the variance welelit matrix for each input pixel.	It is this “instrument Jones matrix”, that is used to form the variance weight matrix for each input pixel.
As we are not accounting for the different primary beam shape. we are not forming the best estimate of the array respouse ina given pixel directiou. it is this simplification that makes the problem tractable in the 32T domain.	As we are not accounting for the different primary beam shape, we are not forming the best estimate of the array response in a given pixel direction, it is this simplification that makes the problem tractable in the 32T domain.
But it should be noted that the weighting scheme will evolve	But it should be noted that the weighting scheme will evolve
To build a sample for flare analvsis. the ANCIIORS data were filtered by counts and variability.	To build a sample for flare analysis, the ANCHORS data were filtered by counts and variability.
The process of selecting sources [or flare analvsis began in June 2008. which means (hat observations must have been completed by June 2007 (o be included in this study.	The process of selecting sources for flare analysis began in June 2008, which means that observations must have been completed by June 2007 to be included in this study.
We find that sources with high Gregorv-Loredo variability grades tend to show more Irequentlv the archetypal flaring behavior of impulsive rise in X-ray luminosity [ollowed by eradual decay to characteristic levels.	We find that sources with high Gregory-Loredo variability grades tend to show more frequently the archetypal flaring behavior of impulsive rise in X-ray luminosity followed by gradual decay to characteristic levels.
A variability grade greater than or equal to 8 (2 confidence level) was the first selection applied to the data.	A variability grade greater than or equal to 8 $>$ confidence level) was the first selection applied to the data.
This first selection vielded about 2000 sources.	This first selection yielded about 2000 sources.
To ensure (hie extraction of a equality spectrum. from these sources we chose observations with at least 1000 counts per source. of which at least 200 come Irom the peak flare segment of the light curve.	To ensure the extraction of a quality spectrum, from these sources we chose observations with at least 1000 counts per source, of which at least 200 come from the peak flare segment of the light curve.
We found this count level to be a good balance between inclusion of the maximum nmunmber of stars for this study and a minimum number of counts needed to extract a quality spectrum for each segment of the light curve.	We found this count level to be a good balance between inclusion of the maximum number of stars for this study and a minimum number of counts needed to extract a quality spectrum for each segment of the light curve.
As a further cut. all sources whose light curve did not show an archetypal [Iare's [ast rise ancl slow decay were discarded for the purposes of this analysis.	As a further cut, all sources whose light curve did not show an archetypal flare's fast rise and slow decay were discarded for the purposes of this analysis.
A light eurve with a slow rise or no decay could indicate either a flare caused by different underlving mechanisms ancl (ius not compatible wilh the procedure developed in 83. or altogether different (vpes of variability. e.g.. rotational modulation.	A light curve with a slow rise or no decay could indicate either a flare caused by different underlying mechanisms and thus not compatible with the procedure developed in 3, or altogether different types of variability, e.g., rotational modulation.
These last selections reduced the data set to 29 stars meeting all criteria [or counts and variability.	These last selections reduced the data set to 29 stars meeting all criteria for counts and variability.
Finally. all sources were checked for photon pile-up.	Finally, all sources were checked for photon pile-up.
For each source. the highest count rate per pixel alter correcting for ofl-axis PSF is compared against the expected pile-up rate.	For each source, the highest count rate per pixel after correcting for off-axis PSF is compared against the expected pile-up rate.
None of the sources selected [or further study suffered [rom photon pile-up greater than5%.	None of the sources selected for further study suffered from photon pile-up greater than.
. Table 1 lists the X-ray properties of the 29 non-COUP stars involved in this study.	Table 1 lists the X-ray properties of the 29 non-COUP stars involved in this study.
The 17 non-COUP observations used in this study are summarized in Table 2.	The 17 non-COUP observations used in this study are summarized in Table 2.
While cluster membership is not explicitly determined in (his study. the existence of 2MASS counterparts for all X-ray sources included and the filtering out of very solt or hard X-ray. photons limit the likelihood of contamination by non-stellar objects for sources in our sample.	While cluster membership is not explicitly determined in this study, the existence of 2MASS counterparts for all X-ray sources included and the filtering out of very soft or hard X-ray photons limit the likelihood of contamination by non-stellar objects for sources in our sample.
The archetvpal fast rise and slow decay. [laring behavior in the sources X-ray light curves is strong additional evidence lor their vouth.	The archetypal fast rise and slow decay flaring behavior in the sources' X-ray light curves is strong additional evidence for their youth.
Cluster membership as displaved in Table 2 is then very. hiehlv probable.	Cluster membership as displayed in Table 2 is then very highly probable.
The selections (hat culiinated in Tables \|. and 2 do not generate an unbiased sample: they specifically select only the brightest flares because only the brightest and longest-Iasting flares will have (he statistics necessary for (nme resolved spectral analysis.	The selections that culminated in Tables 1 and 2 do not generate an unbiased sample: they specifically select only the brightest flares because only the brightest and longest-lasting flares will have the statistics necessary for time resolved spectral analysis.
Moreover. because huninositwv scales with the square of distance. flares from more distant. clusters are more intvinsically intense than (hose in nearby clusters.	Moreover, because luminosity scales with the square of distance, flares from more distant clusters are more intrinsically intense than those in nearby clusters.
For example. flares on sources in M 17 (2100 pe away) must about LOO times more luminous than those in the Serpeus Cloud Core (260 pe away: though perhaps a bit farther. see Winston 2010) to be time-resolved with	For example, flares on sources in M 17 (2100 pc away) must about 100 times more luminous than those in the Serpens Cloud Core (260 pc away; though perhaps a bit further, see Winston 2010) to be time-resolved with
e The remnant euvelope shrinks back onto the white dwarf surface. reduci its moment of inertia aud thus decreasing the spin period to 0.137d.	$\bullet$ The remnant envelope shrinks back onto the white dwarf surface, reducing its moment of inertia and thus decreasing the spin period to 0.137d.
Ciuvenutly. therefore. V1500 Cre displavs two periods.	Currently, therefore, V1500 Cyg displays two periods.
There is a polariuctric signal at 0.127d which is the spin period of the white dwarf. and a photometric signal at 0.1IU which is the orbital period of the Jinary.	There is a polarimetric signal at 0.137d which is the spin period of the white dwarf, and a photometric signal at 0.140d which is the orbital period of the binary.
Aside from Sclunidt et als spectroplotometiy outlined in Section 1.. further evidence that the secondary star iow dominates the photometric modulation comes fron ιο fact hat he timing of flux maxinuni in our own clata natches he orbital ephemeris (sec Section L)).	Aside from Schmidt et al's spectrophotometry outlined in Section \ref{intro}, further evidence that the secondary star now dominates the photometric modulation comes from the fact that the timing of flux maximum in our own data matches the orbital ephemeris (see Section \ref{5:obs}) ).
It is also clear however. frou the presence of flickering aud ie slightly asyaunietrical B baud leht-curves. that there Is sole contamination of the orbital modulation.	It is also clear however, from the presence of flickering and the slightly asymmetrical B band light-curves, that there is some contamination of the orbital modulation.
This could either be due to the accretion stream. or perhaps )ecause the spin period of the white dwarf may also have a photometric signal (Pavlenko Pelt 1988).	This could either be due to the accretion stream, or perhaps because the spin period of the white dwarf may also have a photometric signal (Pavlenko Pelt 1988).
Prialmik (1986) modelled the evolution of a classical nova through a complete cycle: accretion. outburst. mass loss. decline aud restuned accretion.	Prialnik (1986) modelled the evolution of a classical nova through a complete cycle; accretion, outburst, mass loss, decline and resumed accretion.
The model is for a 1.25.U... C-O white dwart.	The model is for a $M_{\odot}$ C-O white dwarf.
The resulting outburst is fast and similar to that observed in V1500 Cre. matching the composition of the ejected euvelope very well.	The resulting outburst is fast and similar to that observed in V1500 Cyg, matching the composition of the ejected envelope very well.
The white dwarf is modelled allowing for heat trausfer via radiation. conduction (Ibeu 1975. I&ovetz Shaviv 1973) aud convection (Milalas 1978).	The white dwarf is modelled allowing for heat transfer via radiation, conduction (Iben 1975, Kovetz Shaviv 1973) and convection (Mihalas 1978).
The result may be fitted well with a power law cooling curve ofthe form where Lois the Iuninositv of the white dwarf aud f is the time since outburst.	The result may be fitted well with a power law cooling curve of the form where L is the luminosity of the white dwarf and $t$ is the time since outburst.
To extend the baseline of photometric amplitude decline we obtained one orbital evele of D aud V baud photometry using the JIKT on La Palma ou the night of 1995 October 3.	To extend the baseline of photometric amplitude decline we obtained one orbital cycle of B and V band photometry using the JKT on La Palma on the night of 1995 October 3.
The TEKI CCD was used with pixels binned 2 by 2 to achieve 0.60.6" pixels iu rapid readout mode.	The TEK4 CCD was used with pixels binned 2 by 2 to achieve 0.6"x0.6" pixels in rapid readout mode.
The seeimg was around 2.0 arcsec.	The seeing was around 2.0 arcsec.
The exposures were typically of 120 secouds with filters being alternated between observations.	The exposures were typically of 120 seconds with filters being alternated between observations.
A bias frame was subtracted from cach image. aud it was then flatfielded using au image of the twilight sky.	A bias frame was subtracted from each image, and it was then flatfielded using an image of the twilight sky.
The counts for V1500 Cre aud various other stars were extracted from cach frame using the optimal weighting procedure deseribed in Navlor (1998).	The counts for V1500 Cyg and various other stars were extracted from each frame using the optimal weighting procedure described in Naylor (1998).
We divided the counts for VI500 Cye by those for star Cl of Ialuzux Semeniik (1987). allowing us to put the lishteurves iu Figures D and 2. onto a magnitude scale.	We divided the counts for V1500 Cyg by those for star C1 of Kaluzny Semeniuk (1987), allowing us to put the lightcurves in Figures \ref{fig:bband} and \ref{fig:vband} onto a magnitude scale.
The times of maxiunun and minim aeree. within errors. with the ephemeris of Semieuiuk. Olech Nalezvty (1995).	The times of maximum and minimum agree, within errors, with the ephemeris of Semeniuk, Olech Nalezyty (1995).
The data do not show a pure heating modulation since there is some evidence of a dip in the helt curve at time 999LOL.	The data do not show a pure heating modulation since there is some evidence of a dip in the light curve at time 994.61.
Such dips are not uuconunuon. see for exalmple the ligbit-curves of Ikaluzux Semmens (1987) whose data are from 10 vears earlier than our observations.	Such dips are not uncommon, see for example the light-curves of Kaluzny Semeniuk (1987) whose data are from 10 years earlier than our observations.
These ireeularities add errors in the estimation of the amplitude of the modulation (see Section 5)).	These irregularities add errors in the estimation of the amplitude of the modulation (see Section \ref{5:Ampvtime}) ).
By searching through various sources a record of the B xuid photometric behaviour observed in VI500 Cre since outburst has been assembled.	By searching through various sources a record of the B band photometric behaviour observed in V1500 Cyg since outburst has been assembled.
Table 3.1 eves a list of all references used.	Table 3.1 gives a list of all references used.
The amplitude of the B baud photometric nodulation versus the time since outburst is plotted iu Figure 3Pa on logarithmic scales.	The amplitude of the B band photometric modulation versus the time since outburst is plotted in Figure \ref{fig:ampVd} on logarithmic scales.
We mauucediatolv note that hese poiuts lie roughly ina straight line. implying a power	We immediately note that these points lie roughly in a straight line, implying a power
around 250350 Myr (Diraffe et al.	around 250–350 Myr (Baraffe et al.
2003) or ~ 260 My (Burrows et al.	2003) or $\sim$ 260 Myr (Burrows et al.
1997) from Imuinosity coustraiuts. oy & 225 Myr (Chabrier et al.	1997) from luminosity constraints, or $\lesssim$ 225 Myr (Chabrier et al.
2000) and 250280 Myr (Daraffe et al.	2000) and 250–280 Myr (Baraffe et al.
2003) from temperature constrains. the dviuuical ifetime of the syste is significantly smaller. = 150 My.	2003) from temperature constrains, the dynamical lifetime of the system is significantly smaller, $\lesssim$ 150 Myr.
Dased ou dviiuical considerations we therefore couchide hat the large companion masses considered in this paper are unlikely aud that the dvuamical state of the svstei. avors ages of ~ LOO Myr ancl companion masses slielitlv low the brow dwarf reeiue.	Based on dynamical considerations we therefore conclude that the large companion masses considered in this paper are unlikely and that the dynamical state of the system favors ages of $\sim$ 100 Myr and companion masses slightly below the brown dwarf regime.
This conclusion is supported by the strong emission of he debris disk at 21 pau. behavior that is very uncomumnion around stars older tha l 500 Αα (Caspar et al.	This conclusion is supported by the strong emission of the debris disk at 24 $\mu$ m, behavior that is very uncommon around stars older than $\sim$ 500 Myr (Gaspar et al.
2009).	2009).
Iu addition. the average temperature of 7350 I& aud the πιοτν of 5 L.. from Sadakane (2006) aud ova et al. (	In addition, the average temperature of 7350 K and the luminosity of 5 $_\odot$ from Sadakane (2006) and Moya et al. (
20105) places the star at or below the zero age main sequence and at a very voung age compared with other A Boo stars (Paunzen et al.	2010b) places the star at or below the zero age main sequence and at a very young age compared with other $\lambda$ Boo stars (Paunzen et al.
2002).	2002).
However. these results are not consistent with the YOSSIijlitv that the system in ~ 1 Cyr old.	However, these results are not consistent with the possibility that the system is $\sim$ 1 Gyr old.
The astroscisinoloey modeliug that suggests this age (Mova et al.	The astroseismology modeling that suggests this age (Moya et al.
2010a) depends critically ou the rotation speed of he star and heuce on the assigned inclination.	2010a) depends critically on the rotation speed of the star and hence on the assigned inclination.
It is oulv valid if the inclination of the equatorial plane of the star is D. 236°.	It is only valid if the inclination of the equatorial plane of the star is $I_* \ge $ $\degr$.
The case for IΞ36° results in two sets of xossilble solutions. a stellar mass of 1.32.1.33 ME. and au aee of 11231625 Mor. or a stellar mass of LLL1.15 M. and aud age of 261380 Myr (with of the models vine iu the generally adopted age range of 30.160 My).	The case for $I_* = 36 \degr$ results in two sets of possible solutions, a stellar mass of 1.32–1.33 $_{\odot}$ and an age of 1123–1625 Myr, or a stellar mass of 1.44–1.45 $_{\odot}$ and and age of 26–430 Myr (with of the models lying in the generally adopted age range of 30–160 Myr).
For £507. Mova et al. (	For $I_* = 50\degr$, Moya et al. (
201048) obtain a stellar mass of 1.32 M... and aud age of 11261186 My.	2010a) obtain a stellar mass of 1.32 $_{\odot}$ and and age of 1126–1486 Myr.
Tlowever. Lafreniere et al. (	However, Lafreniere et al. (
2009) suggestedOO that the inchnation of the orbital plane of the plaucts is fy. = frou astrometry observations with a 10 vear baseline (using NICALOS archival data).	2009) suggested that the inclination of the orbital plane of the planets is $I_{\rm pl}$ = $\degr$ from astrometry observations with a 10 year baseline (using NICMOS archival data).
Su et al. (	Su et al. (
2009) reported roni the degree of sviuncetry of the spatially resolved. 70 jaa disk that the iucliuatiou of the disk is cousisteut with reine face-on and uulikelv to be Zi7 25°.	2009) reported from the degree of symmetry of the spatially resolved 70 $\mu$ m disk that the inclination of the disk is consistent with being face-on and unlikely to be $I_{\rm disk} >$ $\degr$.
We have re-evaluated these observations aud place a 26 upper μπιτ o fü of 107.	We have re-evaluated these observations and place a $\sigma$ upper limit to $I_{\rm disk}$ of $\degr$.
Upcoming JCMT/SCUBA-2. Herschel and UST nuagiuse of IIR 5799 will be able to further constrain the inclination of the debris disk.	Upcoming JCMT/SCUBA-2, Herschel and HST imaging of HR 8799 will be able to further constrain the inclination of the debris disk.
If we assume ji the equatorial plane of the star coincides with 16 orbital plane of the planets aud the dust-produciug lanetesinals Ge. Le=Ly fug. these observational constraints indicate that iudeed the inchnation of the star is likely inside the 18 7«Z£.36" range where the astroseismologv models cannot be applied. or that it is lose to 36° where they permit solutions cousisteut with +je c 100 Myr age derived from other considerations.	If we assume that the equatorial plane of the star coincides with the orbital plane of the planets and the dust-producing planetesimals (i.e. $I_{} = I_{\rm pl} = I_{\rm disk}$ ), these observational constraints indicate that indeed the inclination of the star is likely inside the 18 $\degr < I_ < 36 \degr$ range where the astroseismology models cannot be applied, or that it is close to $\degr$ where they permit solutions consistent with the $\sim$ 100 Myr age derived from other considerations.
We conclude that TR 8799 is voune. with au age of ~ 100 Myr.	We conclude that HR 8799 is young, with an age of $\sim$ 100 Myr.
Its companious are therefore likely all to be very lnassive planets. just below the mass laut for deuterim burniug.	Its companions are therefore likely all to be very massive planets, just below the mass limit for deuterium burning.
We thank Πα. Levisou for providing skecl-Syv\IBA for the cdyvuamical simulations and AL Janson. D. Fabrvckv. R. Murrav-Cli. A. Mova aud P. Isalas for useful discussion.	We thank Hal Levison for providing skeel-SyMBA for the dynamical simulations and M. Janson, D. Fabrycky, R. Murray-Clay, A. Moya and P. Kalas for useful discussion.
ΑλΕλ acknowledees fuudius from the Spanish MICTNN (Ramoun y Cajal Program RYC- and erants AYA2009-07301L and Cousolider lugeuio 20L0CSD2009-000385.	A.M.M. acknowledges funding from the Spanish MICINN (Ramónn y Cajal Program RYC-2007-00612, and grants AYA2009-07304 and Consolider Ingenio 2010CSD2009-00038).
of Cambridge. Case Western Reserve University. University of Chicago. Drexel University. Fermilab. the Institute for Advanced Study. the Japan Participation Group. Johns Hopkins University. he Joint Institute for Nuclear Astrophysics. the Kavli Institute or Particle Astrophysics and Cosmology. the Korean Scientist Group. the Chinese Academy of Sciences (LAMOST). Los Alamos ational Laboratory. the Max-Planck-Institute for Astronomy (MPIA). the Max-Planck-Institute for Astrophysics (MPA). New Texico State University. Ohio State University. University of Pittsburgh. University of Portsmouth. Princeton University. he United States Naval. Observatory. and the University of Washington.	of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.
Finally, we note that our findings for tail-making during galaxy-galaxy encounters may apply to gas as well.	Finally, we note that our findings for tail-making during galaxy-galaxy encounters may apply to gas as well.
Thus, our calculations be used to understand the of the streams and bridges of stars associated with dwarf galaxies, as recently discoveredmay in the Panda survey of M31 geometry(McConnachiegasetal.2009)..	Thus, our calculations may be used to understand the geometry of the gas streams and bridges of stars associated with dwarf galaxies, as recently discovered in the Panda survey of M31 \citep{McConn09}.
Other potential applications include warps and heating of galactic disks by tidal perturbations, understanding the conditions for long tails to be produced in galaxy collisions depending on the dark matter distribution in the halos, and identifying situations where stars can be unbound in galaxy interactions and be ejected into the intergalactic medium or populate the field of groups and clusters.	Other potential applications include warps and heating of galactic disks by tidal perturbations, understanding the conditions for long tails to be produced in galaxy collisions depending on the dark matter distribution in the halos, and identifying situations where stars can be unbound in galaxy interactions and be ejected into the intergalactic medium or populate the field of groups and clusters.
Moreover, our theory may be relevant to	Moreover, our theory may be relevant to
In relatively nearby starburst galaxies. recent star formation activity. has (vpically been resolved into massive sliw clusters.	In relatively nearby starburst galaxies, recent star formation activity has typically been resolved into massive star clusters.
While old elobular clusters are ubiquitous at the current epoch. only over the past decade have we begun to find their vounger ancl bluer siblings. "super star clusters” (SSC85). in signilicant numbers (Whitmore2000).	While old globular clusters are ubiquitous at the current epoch, only over the past decade have we begun to find their younger and bluer siblings, “super star clusters" (SSCs), in significant numbers \citep{whitmore00}.
. The large number of elobular clusters in the οσα universe appear to have been formed during the early stages of galaxy evolution.	The large number of globular clusters in the local universe appear to have been formed during the early stages of galaxy evolution.

The observed-to-intrinsic luminosity ratio in the 2-10 keV band is then 0.675.! per cent for IRAS 0910444109.

The observed-to-intrinsic luminosity ratio in the 2–10 keV band is then $0.67\gamma^{-1}$ per cent for IRAS 09104+4109.

Tf the ratio of the 2-10 keV to bolometric luminosities is assumed to be {ηνyaar=0.05fixepaotoos-. typical for quasars (e.g. Elvis et al 1994). the bolometric luminosity is +.

If the ratio of the 2–10 keV to bolometric luminosities is assumed to be $f_{\rm HX/Bol}=0.05f_{\rm HX/Bol,0.05}$, typical for quasars (e.g., Elvis et al 1994), the bolometric luminosity is $1.5\times 10^{46}f_{\rm HX/Bol,0.05}^{-1} \gamma h^{-2}$ .

This is comparable to Li,=2.3 . and consistent. with. a picture. where the bulk of. the infrared emission is powered by a buried quasar through dust reradiation.

This is comparable to $L_{\rm ir}=2.3\times 10^{46}h^{-2}$ , and consistent with a picture where the bulk of the infrared emission is powered by a buried quasar through dust reradiation.

The radio source has an intermediate power between

measurements of varying fidelity as a fuuetion of hour angle. a feature that must be incorporated in the combination of suapshots by a suitable welghtiug scheme.

measurements of varying fidelity as a function of hour angle, a feature that must be incorporated in the combination of snapshots by a suitable weighting scheme.

For each output image pixel we are attempting to form the best estimate of the sky polarized brightuess (by) from the combination of. η individual measurements (bjj.;). each of different weight (105:,:;).

For each output image pixel we are attempting to form the best estimate of the sky polarized brightness $\bar{\bf b}_{pix}$ ) from the combination of, $n$ individual measurements ${\bf b}_{pix,j}$ ), each of different weight $w_{pix,j}$ ).

The best such combined estimate is the weighted mean: where the weights (αρ) are generally the reciprocal of the pixel variance.

The best such combined estimate is the weighted mean: where the weights $w_{pix,j}$ ) are generally the reciprocal of the pixel variance.

The inverse-variauce of the MWA image pixels is best represented by the direction dependent polarimetric tile response iu the cirectiou of that pixel (the power respouse).

The inverse-variance of the MWA image pixels is best represented by the direction dependent polarimetric tile response in the direction of that pixel (the power response).

We are weightiug therelore by the power signal to noise ratio. uicer the assumption that the the thermal uoise is direction indepeudent.

We are weighting therefore by the power signal to noise ratio, under the assumption that the the thermal noise is direction independent.

Following the formalisin developed iu the series of papers: 2: 2: 2: ?. aud ο. the brightuess of au image pixel iu the measured. or instrument [rate is: where b are the coherency Lvectors. the superscript m iucicates “measured” aud should be taken to imply au instrument polarization basis (p. q) aud 5 represents a Stokes (4.Q.U. V) basis.

Following the formalism developed in the series of papers: \cite{hbs96}; \cite{shb96}; \cite{hb96}; \cite{ham00} and \cite{ham06}, the brightness of an image pixel in the measured, or instrument frame is: where ${\bf b}$ are the coherency 4-vectors, the superscript ${\bf m}$ indicates “measured” and should be taken to imply an instrument polarization basis $p,q$ ) and ${\bf s}$ represents a Stokes $I, Q, U, V$ ) basis.

The matrix 8 trausforms polarization from the instrument basis (orthogonal linear feeds) to the Stokes basis ou the sky aud is The MWA-32T is an interferometer comprising 32 dillerent elements but. as is described in §2.. each visibility is calibrated by au equalization process that allows us to approximate the iustruimenut polarimetric response by anmatrix.

The matrix $\stokes$ transforms polarization from the instrument basis (orthogonal linear feeds) to the Stokes basis on the sky and is The MWA-32T is an interferometer comprising 32 different elements but, as is described in \ref{sec:calibration}, each visibility is calibrated by an equalization process that allows us to approximate the instrument polarimetric response by an.

The 2x2 instrument Joues matrix J. determineW. how the generally non-orthiogonal projected receptors (p.q) respoud to the iucideut orthogoual sky polarization (X. Y).

The 2x2 instrument Jones matrix $\jones$, determines how the generally non-orthogonal projected receptors ) respond to the incident orthogonal sky polarization (X, Y).

As the projection of the tile receptors changes as a function of position the Jones matrix is different [or every pixel.

As the projection of the tile receptors changes as a function of position the Jones matrix is different for every pixel.

[tis this “instrument Joues matrix”. that is used to form the variance welelit matrix for each input pixel.

It is this “instrument Jones matrix”, that is used to form the variance weight matrix for each input pixel.

As we are not accounting for the different primary beam shape. we are not forming the best estimate of the array respouse ina given pixel directiou. it is this simplification that makes the problem tractable in the 32T domain.

As we are not accounting for the different primary beam shape, we are not forming the best estimate of the array response in a given pixel direction, it is this simplification that makes the problem tractable in the 32T domain.

But it should be noted that the weighting scheme will evolve

To build a sample for flare analvsis. the ANCIIORS data were filtered by counts and variability.

To build a sample for flare analysis, the ANCHORS data were filtered by counts and variability.

The process of selecting sources [or flare analvsis began in June 2008. which means (hat observations must have been completed by June 2007 (o be included in this study.

The process of selecting sources for flare analysis began in June 2008, which means that observations must have been completed by June 2007 to be included in this study.

We find that sources with high Gregorv-Loredo variability grades tend to show more Irequentlv the archetypal flaring behavior of impulsive rise in X-ray luminosity [ollowed by eradual decay to characteristic levels.

We find that sources with high Gregory-Loredo variability grades tend to show more frequently the archetypal flaring behavior of impulsive rise in X-ray luminosity followed by gradual decay to characteristic levels.

A variability grade greater than or equal to 8 (2 confidence level) was the first selection applied to the data.

A variability grade greater than or equal to 8 $>$ confidence level) was the first selection applied to the data.

This first selection vielded about 2000 sources.

This first selection yielded about 2000 sources.

To ensure (hie extraction of a equality spectrum. from these sources we chose observations with at least 1000 counts per source. of which at least 200 come Irom the peak flare segment of the light curve.

To ensure the extraction of a quality spectrum, from these sources we chose observations with at least 1000 counts per source, of which at least 200 come from the peak flare segment of the light curve.

We found this count level to be a good balance between inclusion of the maximum nmunmber of stars for this study and a minimum number of counts needed to extract a quality spectrum for each segment of the light curve.

We found this count level to be a good balance between inclusion of the maximum number of stars for this study and a minimum number of counts needed to extract a quality spectrum for each segment of the light curve.

As a further cut. all sources whose light curve did not show an archetypal [Iare's [ast rise ancl slow decay were discarded for the purposes of this analysis.

As a further cut, all sources whose light curve did not show an archetypal flare's fast rise and slow decay were discarded for the purposes of this analysis.

A light eurve with a slow rise or no decay could indicate either a flare caused by different underlving mechanisms ancl (ius not compatible wilh the procedure developed in 83. or altogether different (vpes of variability. e.g.. rotational modulation.

A light curve with a slow rise or no decay could indicate either a flare caused by different underlying mechanisms and thus not compatible with the procedure developed in 3, or altogether different types of variability, e.g., rotational modulation.

These last selections reduced the data set to 29 stars meeting all criteria [or counts and variability.

These last selections reduced the data set to 29 stars meeting all criteria for counts and variability.

Finally. all sources were checked for photon pile-up.

Finally, all sources were checked for photon pile-up.

For each source. the highest count rate per pixel alter correcting for ofl-axis PSF is compared against the expected pile-up rate.

For each source, the highest count rate per pixel after correcting for off-axis PSF is compared against the expected pile-up rate.

None of the sources selected [or further study suffered [rom photon pile-up greater than5%.

None of the sources selected for further study suffered from photon pile-up greater than.

. Table 1 lists the X-ray properties of the 29 non-COUP stars involved in this study.

Table 1 lists the X-ray properties of the 29 non-COUP stars involved in this study.

The 17 non-COUP observations used in this study are summarized in Table 2.

While cluster membership is not explicitly determined in (his study. the existence of 2MASS counterparts for all X-ray sources included and the filtering out of very solt or hard X-ray. photons limit the likelihood of contamination by non-stellar objects for sources in our sample.

While cluster membership is not explicitly determined in this study, the existence of 2MASS counterparts for all X-ray sources included and the filtering out of very soft or hard X-ray photons limit the likelihood of contamination by non-stellar objects for sources in our sample.

The archetvpal fast rise and slow decay. [laring behavior in the sources X-ray light curves is strong additional evidence lor their vouth.

The archetypal fast rise and slow decay flaring behavior in the sources' X-ray light curves is strong additional evidence for their youth.

Cluster membership as displaved in Table 2 is then very. hiehlv probable.

Cluster membership as displayed in Table 2 is then very highly probable.

The selections (hat culiinated in Tables |. and 2 do not generate an unbiased sample: they specifically select only the brightest flares because only the brightest and longest-Iasting flares will have (he statistics necessary for (nme resolved spectral analysis.

The selections that culminated in Tables 1 and 2 do not generate an unbiased sample: they specifically select only the brightest flares because only the brightest and longest-lasting flares will have the statistics necessary for time resolved spectral analysis.

Moreover. because huninositwv scales with the square of distance. flares from more distant. clusters are more intvinsically intense than (hose in nearby clusters.

Moreover, because luminosity scales with the square of distance, flares from more distant clusters are more intrinsically intense than those in nearby clusters.

For example. flares on sources in M 17 (2100 pe away) must about LOO times more luminous than those in the Serpeus Cloud Core (260 pe away: though perhaps a bit farther. see Winston 2010) to be time-resolved with

For example, flares on sources in M 17 (2100 pc away) must about 100 times more luminous than those in the Serpens Cloud Core (260 pc away; though perhaps a bit further, see Winston 2010) to be time-resolved with

e The remnant euvelope shrinks back onto the white dwarf surface. reduci its moment of inertia aud thus decreasing the spin period to 0.137d.

$\bullet$ The remnant envelope shrinks back onto the white dwarf surface, reducing its moment of inertia and thus decreasing the spin period to 0.137d.

Ciuvenutly. therefore. V1500 Cre displavs two periods.

Currently, therefore, V1500 Cyg displays two periods.

There is a polariuctric signal at 0.127d which is the spin period of the white dwarf. and a photometric signal at 0.1IU which is the orbital period of the Jinary.

There is a polarimetric signal at 0.137d which is the spin period of the white dwarf, and a photometric signal at 0.140d which is the orbital period of the binary.

Aside from Sclunidt et als spectroplotometiy outlined in Section 1.. further evidence that the secondary star iow dominates the photometric modulation comes fron ιο fact hat he timing of flux maxinuni in our own clata natches he orbital ephemeris (sec Section L)).

Aside from Schmidt et al's spectrophotometry outlined in Section \ref{intro}, further evidence that the secondary star now dominates the photometric modulation comes from the fact that the timing of flux maximum in our own data matches the orbital ephemeris (see Section \ref{5:obs}) ).

It is also clear however. frou the presence of flickering aud ie slightly asyaunietrical B baud leht-curves. that there Is sole contamination of the orbital modulation.

It is also clear however, from the presence of flickering and the slightly asymmetrical B band light-curves, that there is some contamination of the orbital modulation.

This could either be due to the accretion stream. or perhaps )ecause the spin period of the white dwarf may also have a photometric signal (Pavlenko Pelt 1988).

This could either be due to the accretion stream, or perhaps because the spin period of the white dwarf may also have a photometric signal (Pavlenko Pelt 1988).

Prialmik (1986) modelled the evolution of a classical nova through a complete cycle: accretion. outburst. mass loss. decline aud restuned accretion.

Prialnik (1986) modelled the evolution of a classical nova through a complete cycle; accretion, outburst, mass loss, decline and resumed accretion.

The model is for a 1.25.U... C-O white dwart.

The model is for a $M_{\odot}$ C-O white dwarf.

The resulting outburst is fast and similar to that observed in V1500 Cre. matching the composition of the ejected euvelope very well.

The resulting outburst is fast and similar to that observed in V1500 Cyg, matching the composition of the ejected envelope very well.

The white dwarf is modelled allowing for heat trausfer via radiation. conduction (Ibeu 1975. I&ovetz Shaviv 1973) aud convection (Milalas 1978).

The white dwarf is modelled allowing for heat transfer via radiation, conduction (Iben 1975, Kovetz Shaviv 1973) and convection (Mihalas 1978).

The result may be fitted well with a power law cooling curve ofthe form where Lois the Iuninositv of the white dwarf aud f is the time since outburst.

The result may be fitted well with a power law cooling curve of the form where L is the luminosity of the white dwarf and $t$ is the time since outburst.

To extend the baseline of photometric amplitude decline we obtained one orbital evele of D aud V baud photometry using the JIKT on La Palma ou the night of 1995 October 3.

To extend the baseline of photometric amplitude decline we obtained one orbital cycle of B and V band photometry using the JKT on La Palma on the night of 1995 October 3.

The TEKI CCD was used with pixels binned 2 by 2 to achieve 0.60.6" pixels iu rapid readout mode.

The TEK4 CCD was used with pixels binned 2 by 2 to achieve 0.6"x0.6" pixels in rapid readout mode.

The seeimg was around 2.0 arcsec.

The seeing was around 2.0 arcsec.

The exposures were typically of 120 secouds with filters being alternated between observations.

The exposures were typically of 120 seconds with filters being alternated between observations.

A bias frame was subtracted from cach image. aud it was then flatfielded using au image of the twilight sky.

A bias frame was subtracted from each image, and it was then flatfielded using an image of the twilight sky.

The counts for V1500 Cre aud various other stars were extracted from cach frame using the optimal weighting procedure deseribed in Navlor (1998).

The counts for V1500 Cyg and various other stars were extracted from each frame using the optimal weighting procedure described in Naylor (1998).

We divided the counts for VI500 Cye by those for star Cl of Ialuzux Semeniik (1987). allowing us to put the lishteurves iu Figures D and 2. onto a magnitude scale.

We divided the counts for V1500 Cyg by those for star C1 of Kaluzny Semeniuk (1987), allowing us to put the lightcurves in Figures \ref{fig:bband} and \ref{fig:vband} onto a magnitude scale.

The times of maxiunun and minim aeree. within errors. with the ephemeris of Semieuiuk. Olech Nalezvty (1995).

The times of maximum and minimum agree, within errors, with the ephemeris of Semeniuk, Olech Nalezyty (1995).

The data do not show a pure heating modulation since there is some evidence of a dip in the helt curve at time 999LOL.

The data do not show a pure heating modulation since there is some evidence of a dip in the light curve at time 994.61.

Such dips are not uuconunuon. see for exalmple the ligbit-curves of Ikaluzux Semmens (1987) whose data are from 10 vears earlier than our observations.

Such dips are not uncommon, see for example the light-curves of Kaluzny Semeniuk (1987) whose data are from 10 years earlier than our observations.

These ireeularities add errors in the estimation of the amplitude of the modulation (see Section 5)).

These irregularities add errors in the estimation of the amplitude of the modulation (see Section \ref{5:Ampvtime}) ).

By searching through various sources a record of the B xuid photometric behaviour observed in VI500 Cre since outburst has been assembled.

By searching through various sources a record of the B band photometric behaviour observed in V1500 Cyg since outburst has been assembled.

Table 3.1 eves a list of all references used.

Table 3.1 gives a list of all references used.

The amplitude of the B baud photometric nodulation versus the time since outburst is plotted iu Figure 3Pa on logarithmic scales.

The amplitude of the B band photometric modulation versus the time since outburst is plotted in Figure \ref{fig:ampVd} on logarithmic scales.

We mauucediatolv note that hese poiuts lie roughly ina straight line. implying a power

We immediately note that these points lie roughly in a straight line, implying a power

around 250350 Myr (Diraffe et al.

around 250–350 Myr (Baraffe et al.

2003) or ~ 260 My (Burrows et al.

2003) or $\sim$ 260 Myr (Burrows et al.

1997) from Imuinosity coustraiuts. oy & 225 Myr (Chabrier et al.

1997) from luminosity constraints, or $\lesssim$ 225 Myr (Chabrier et al.

2000) and 250280 Myr (Daraffe et al.

2000) and 250–280 Myr (Baraffe et al.

2003) from temperature constrains. the dviuuical ifetime of the syste is significantly smaller. = 150 My.

2003) from temperature constrains, the dynamical lifetime of the system is significantly smaller, $\lesssim$ 150 Myr.

Dased ou dviiuical considerations we therefore couchide hat the large companion masses considered in this paper are unlikely aud that the dvuamical state of the svstei. avors ages of ~ LOO Myr ancl companion masses slielitlv low the brow dwarf reeiue.

Based on dynamical considerations we therefore conclude that the large companion masses considered in this paper are unlikely and that the dynamical state of the system favors ages of $\sim$ 100 Myr and companion masses slightly below the brown dwarf regime.

This conclusion is supported by the strong emission of he debris disk at 21 pau. behavior that is very uncomumnion around stars older tha l 500 Αα (Caspar et al.

This conclusion is supported by the strong emission of the debris disk at 24 $\mu$ m, behavior that is very uncommon around stars older than $\sim$ 500 Myr (Gaspar et al.

2009).

Iu addition. the average temperature of 7350 I& aud the πιοτν of 5 L.. from Sadakane (2006) aud ova et al. (

In addition, the average temperature of 7350 K and the luminosity of 5 $_\odot$ from Sadakane (2006) and Moya et al. (

20105) places the star at or below the zero age main sequence and at a very voung age compared with other A Boo stars (Paunzen et al.

2010b) places the star at or below the zero age main sequence and at a very young age compared with other $\lambda$ Boo stars (Paunzen et al.

2002).

However. these results are not consistent with the YOSSIijlitv that the system in ~ 1 Cyr old.

However, these results are not consistent with the possibility that the system is $\sim$ 1 Gyr old.

The astroscisinoloey modeliug that suggests this age (Mova et al.

The astroseismology modeling that suggests this age (Moya et al.

2010a) depends critically ou the rotation speed of he star and heuce on the assigned inclination.

2010a) depends critically on the rotation speed of the star and hence on the assigned inclination.

It is oulv valid if the inclination of the equatorial plane of the star is D. 236°.

It is only valid if the inclination of the equatorial plane of the star is $I_* \ge $ $\degr$.

The case for IΞ36° results in two sets of xossilble solutions. a stellar mass of 1.32.1.33 ME. and au aee of 11231625 Mor. or a stellar mass of LLL1.15 M. and aud age of 261380 Myr (with of the models vine iu the generally adopted age range of 30.160 My).

The case for $I_* = 36 \degr$ results in two sets of possible solutions, a stellar mass of 1.32–1.33 $_{\odot}$ and an age of 1123–1625 Myr, or a stellar mass of 1.44–1.45 $_{\odot}$ and and age of 26–430 Myr (with of the models lying in the generally adopted age range of 30–160 Myr).

For £507. Mova et al. (

For $I_* = 50\degr$, Moya et al. (

201048) obtain a stellar mass of 1.32 M... and aud age of 11261186 My.

2010a) obtain a stellar mass of 1.32 $_{\odot}$ and and age of 1126–1486 Myr.

Tlowever. Lafreniere et al. (

However, Lafreniere et al. (

2009) suggestedOO that the inchnation of the orbital plane of the plaucts is fy. = frou astrometry observations with a 10 vear baseline (using NICALOS archival data).

2009) suggested that the inclination of the orbital plane of the planets is $I_{\rm pl}$ = $\degr$ from astrometry observations with a 10 year baseline (using NICMOS archival data).

Su et al. (

2009) reported roni the degree of sviuncetry of the spatially resolved. 70 jaa disk that the iucliuatiou of the disk is cousisteut with reine face-on and uulikelv to be Zi7 25°.

2009) reported from the degree of symmetry of the spatially resolved 70 $\mu$ m disk that the inclination of the disk is consistent with being face-on and unlikely to be $I_{\rm disk} >$ $\degr$.

We have re-evaluated these observations aud place a 26 upper μπιτ o fü of 107.

We have re-evaluated these observations and place a $\sigma$ upper limit to $I_{\rm disk}$ of $\degr$.

Upcoming JCMT/SCUBA-2. Herschel and UST nuagiuse of IIR 5799 will be able to further constrain the inclination of the debris disk.

Upcoming JCMT/SCUBA-2, Herschel and HST imaging of HR 8799 will be able to further constrain the inclination of the debris disk.

If we assume ji the equatorial plane of the star coincides with 16 orbital plane of the planets aud the dust-produciug lanetesinals Ge. Le=Ly fug. these observational constraints indicate that iudeed the inchnation of the star is likely inside the 18 7«Z£.36" range where the astroseismologv models cannot be applied. or that it is lose to 36° where they permit solutions cousisteut with +je c 100 Myr age derived from other considerations.

If we assume that the equatorial plane of the star coincides with the orbital plane of the planets and the dust-producing planetesimals (i.e. $I_{*} = I_{\rm pl} = I_{\rm disk}$ ), these observational constraints indicate that indeed the inclination of the star is likely inside the 18 $\degr < I_* < 36 \degr$ range where the astroseismology models cannot be applied, or that it is close to $\degr$ where they permit solutions consistent with the $\sim$ 100 Myr age derived from other considerations.

We conclude that TR 8799 is voune. with au age of ~ 100 Myr.

We conclude that HR 8799 is young, with an age of $\sim$ 100 Myr.

Its companious are therefore likely all to be very lnassive planets. just below the mass laut for deuterim burniug.

Its companions are therefore likely all to be very massive planets, just below the mass limit for deuterium burning.

We thank Πα. Levisou for providing skecl-Syv\IBA for the cdyvuamical simulations and AL Janson. D. Fabrvckv. R. Murrav-Cli. A. Mova aud P. Isalas for useful discussion.

We thank Hal Levison for providing skeel-SyMBA for the dynamical simulations and M. Janson, D. Fabrycky, R. Murray-Clay, A. Moya and P. Kalas for useful discussion.

ΑλΕλ acknowledees fuudius from the Spanish MICTNN (Ramoun y Cajal Program RYC- and erants AYA2009-07301L and Cousolider lugeuio 20L0CSD2009-000385.

A.M.M. acknowledges funding from the Spanish MICINN (Ramónn y Cajal Program RYC-2007-00612, and grants AYA2009-07304 and Consolider Ingenio 2010CSD2009-00038).

of Cambridge. Case Western Reserve University. University of Chicago. Drexel University. Fermilab. the Institute for Advanced Study. the Japan Participation Group. Johns Hopkins University. he Joint Institute for Nuclear Astrophysics. the Kavli Institute or Particle Astrophysics and Cosmology. the Korean Scientist Group. the Chinese Academy of Sciences (LAMOST). Los Alamos ational Laboratory. the Max-Planck-Institute for Astronomy (MPIA). the Max-Planck-Institute for Astrophysics (MPA). New Texico State University. Ohio State University. University of Pittsburgh. University of Portsmouth. Princeton University. he United States Naval. Observatory. and the University of Washington.

of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.

Finally, we note that our findings for tail-making during galaxy-galaxy encounters may apply to gas as well.

Thus, our calculations be used to understand the of the streams and bridges of stars associated with dwarf galaxies, as recently discoveredmay in the Panda survey of M31 geometry(McConnachiegasetal.2009)..

Thus, our calculations may be used to understand the geometry of the gas streams and bridges of stars associated with dwarf galaxies, as recently discovered in the Panda survey of M31 \citep{McConn09}.

Other potential applications include warps and heating of galactic disks by tidal perturbations, understanding the conditions for long tails to be produced in galaxy collisions depending on the dark matter distribution in the halos, and identifying situations where stars can be unbound in galaxy interactions and be ejected into the intergalactic medium or populate the field of groups and clusters.

Moreover, our theory may be relevant to

In relatively nearby starburst galaxies. recent star formation activity. has (vpically been resolved into massive sliw clusters.

In relatively nearby starburst galaxies, recent star formation activity has typically been resolved into massive star clusters.

While old elobular clusters are ubiquitous at the current epoch. only over the past decade have we begun to find their vounger ancl bluer siblings. "super star clusters” (SSC85). in signilicant numbers (Whitmore2000).

While old globular clusters are ubiquitous at the current epoch, only over the past decade have we begun to find their younger and bluer siblings, “super star clusters" (SSCs), in significant numbers \citep{whitmore00}.

. The large number of elobular clusters in the οσα universe appear to have been formed during the early stages of galaxy evolution.

The large number of globular clusters in the local universe appear to have been formed during the early stages of galaxy evolution.