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All three show a series of correlated decreases i- flux. | All three show a series of correlated decreases in flux. |
The Oogaps οποσα the successive trends correspond o the trausitious between observing cycles (section £)). | The gaps between the successive trends correspond to the transitions between observing cycles (section \ref{data}) ). |
This effect is caused by the cecentering of the star ou he spectrograph slit duiug a cvele. | This effect is caused by the decentering of the star on the spectrograph slit during a cycle. |
To affect the relative spectra. the degree of deceutering must be different for he two objects. either because the slit did not remai- xuwallel to the Lue joiuing the stars. or because the slit edges were not sufficiently parallel to cach other. | To affect the relative spectra, the degree of decentering must be different for the two objects, either because the slit did not remain parallel to the line joining the stars, or because the slit edges were not sufficiently parallel to each other. |
Althousg[um such deceuteriug was expected (from differeutial flexiug )etwoeen the guide camera aud the instrument). it was nof expected to be as large as 0.2nunags per hour. | Although such decentering was expected (from differential flexing between the guide camera and the instrument), it was not expected to be as large as mags per hour. |
Thus we cannot perform pure differential spectropliotoimietry with hese data. | Thus we cannot perform pure differential spectrophotometry with these data. |
Colour variability within 2MIII5 alone also turus out o be affected by the same problem. | Colour variability within 2M1145 alone also turns out to be affected by the same problem. |
ΤΟΠΟ shows the variation iu the colours Lh2 aud kl fommed from the nominal (Gather than relative) spectra of 2MI11I5 and the reference star. | \\ref{res2} shows the variation in the colours $-$ h2 and $-$ k1 formed from the nominal (rather than relative) spectra of 2M1145 and the reference star. |
We imuuediately see that within each evele both objects appear to eet bluer. nunore rec licht. than blue light is lost from the slit decentering. | We immediately see that within each cycle both objects appear to get bluer, more red light than blue light is lost from the slit decentering. |
This is most likely due to atinosplieric refraction. such that the redder part of the refracted image is that which is more deceutered with respect to the slit. | This is most likely due to atmospheric refraction, such that the redder part of the refracted image is that which is more decentered with respect to the slit. |
Iu spite of these problems. we fud hatreliable. in the sense that there is no systematic trend within a evcle. | In spite of these problems, we find that, in the sense that there is no systematic trend within a cycle. |
This can be seen in Fie. | This can be seen in Fig. |
10. for the colours 19. h2. h2 aud kl. | \ref{res3}
for the colours $-$ j2, $-$ h2, $-$ h2 and $-$ k1. |
That the colours formed roni the relative spectra can be reliable despite he problems scen in rofresl and 9 can be uuderstooc from the fact that both 2MITIS aud the refereuce star show the same degree of "blueuimeg as a result of the slit deceuterug (see Fig. 9)). | That the colours formed from the relative spectra can be reliable despite the problems seen in \\ref{res1} and \ref{res2} can be understood from the fact that both 2M1145 and the reference star show the same degree of `bluening' as a result of the slit decentering (see Fig. \ref{res2}) ). |
It is fainlv clear from a visual inspection of ο that tone of the four light curves shows good evidence for variability. | It is fairly clear from a visual inspection of \\ref{res3} that none of the four light curves shows good evidence for variability. |
This cau be put ou a quantitative basis. using. a statistical⋅⋅ test. for. example the ttest (see section {το ΟΝΕ. | This can be put on a quantitative basis using a statistical test, for example the test (see section 4.1 of BJM). |
The probabilities hat there js a real variability nuot related to the noise) are shown in Table L. | The probabilities that there is a real variability not related to the noise) are shown in Table \ref{amps}. |
This table also gives the amplitudes of the relative colour light curves. | This table also gives the amplitudes of the relative colour light curves. |
It is tempting to study the ratios of these amplitudes aud compare them with the predictious given in section 3.. | It is tempting to study the ratios of these amplitudes and compare them with the predictions given in section \ref{predictions}. |
However. the dominaut contribution to these anplitudes is uoise. so we would first have to statistically subtract the nolse contribution. | However, the dominant contribution to these amplitudes is noise, so we would first have to statistically subtract the noise contribution. |
This | This |
= l]mJy and z & 2 would have silicate absorption or emission features anc no PAII features. | $\ga$ 1 mJy and z $\sim$ 2 would have silicate absorption or emission features and no PAH features. |
Because PAIL sources have been discovered by Spifzer with [,(24jm)) = 1 mJy and z ~ 2. substantial huminositv evolution is required for sources with PAITL-dominated spectra. | Because PAH sources have been discovered by $Spitzer$ with $_{\nu}$ ) $\ga$ 1 mJy and z $\sim$ 2, substantial luminosity evolution is required for sources with PAH-dominated spectra. |
For the pure starbursts. dispersions among PALL strength. [NeLT] emission line strength. and dust continuum strength are illustrated to estimate (he resulting dispersions among star formation rates that would be derived [rom these three independent indicators of SFR. | For the pure starbursts, dispersions among PAH strength, [NeIII] emission line strength, and dust continuum strength are illustrated to estimate the resulting dispersions among star formation rates that would be derived from these three independent indicators of SFR. |
It is found that the dispersions indicate an uncertainty of about a [actor of (wo in deriving the SER from these different indicators. | It is found that the dispersions indicate an uncertainty of about a factor of two in deriving the SFR from these different indicators. |
We thank P. Hall lor help in improving our URS spectral analvsis with SMART. | We thank P. Hall for help in improving our IRS spectral analysis with SMART. |
This work is based primarily on observations made with the Spitzer Space Telescope. which is operated by the Jet Propulsion Laboratory. California Institute of Technology under NASA contract. 1407. | This work is based primarily on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407. |
Support for this work by the IRS GTO team at Cornell University was provided by NASA through Contract Number 1257184 issued by JPL/Caltech. | Support for this work by the IRS GTO team at Cornell University was provided by NASA through Contract Number 1257184 issued by JPL/Caltech. |
This research has macle use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration. | This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. |
dispersion Anauib=111x 27428 lo for stars brighter than ms,δι Μας, assundug a ucelicible effect of binary orbital motions (Cürard ct al. | dispersion $\sigma_{\rm
pm 1D}=141\pm27 \mu$ as $^{-1}$ for stars brighter than $m_{814}\approx$ mag, assuming a negligible effect of binary orbital motions (Girard et al. |
1989). | 1989). |
This corresponds to dey=15 θα | at a distance of kkpe. | This corresponds to $\sigma_{\rm cl
1D}=4.5\pm0.8$ km $^{-1}$ at a distance of kpc. |
The coustaut velocity dispersion for stars in the mass range 9MM.. indicates a lack of equipartition of energv amoug cluster members. | The constant velocity dispersion for stars in the mass range $_\odot$ indicates a lack of equipartition of energy among cluster members. |
This provides a strong indication that NYC is far from virial equilibrium. | This provides a strong indication that NYC is far from virial equilibrium. |
Nevertheless. an upper lait of the cluster mass can be obtained by deriving the virial mass Maga, from the observed velocity dispersionLOST): where 422.5 (weakly depeudius ou cluster deusitv structure). rj is the halfanass radius. σοοο ds the threc-dinensional velocitv dispersion. audG ids the eravitational constant. | Nevertheless, an upper limit of the cluster mass can be obtained by deriving the virial mass ${\rm M}_{\rm dyn}$ from the observed velocity dispersion: where $\eta\approx2.5$ (weakly depending on cluster density structure), $r_{\rm
h}$ is the half-mass radius, $\sigma_{\rm cl 3D}$ is the three-dimensional velocity dispersion, and is the gravitational constant. |
NYC is ass segregated with df$ core radius duereasing with decreasing stellar mass2002). | NYC is mass segregated with its core radius increasing with decreasing stellar mass. |
.. À— lower limit ou ry for the ligh-nass stars as derived from[ST— data is comparable to the core radius of z0.2ppe (Stolte 2003). whereas Taravama et (2008) based on the analvsis of ucarufrared adaptive optics data estimate rj — 1.5ppe for stars in the mass range MM... | A lower limit on $r_{\rm h}$ for the high-mass stars as derived from data is comparable to the core radius of $\approx$ pc (Stolte 2003), whereas Harayama et (2008) based on the analysis of near-infrared adaptive optics data estimate $r_{\rm h}$ = pc for stars in the mass range $_\odot$. |
If woe assume Fc ppc aud a three-dinenusional velocity dispersion of esp=v4«L5E δα 3. we derive Maga,=1760043800NL... | If we assume $r_{\rm h}$ = pc and a three-dimensional velocity dispersion of $\sigma_{\rm cl 3D} =
\sqrt{3} \times 4.5\pm0.8$ km $^{-1}$, we derive $M_{\rm
dyn}=17600\pm3800\,\rm{M}_\odot$. |
Considering that this cinamical mass estimate provides an upper limit. it is in aerecient with photometric studies of NYC. which assigned masses to individual stars. and cstimated the total stellar ass to AL,zs1000016000NL..2008).. | Considering that this dynamical mass estimate provides an upper limit, it is in agreement with photometric studies of NYC, which assigned masses to individual stars, and estimated the total stellar mass to $M_{cl}\approx10000-16000\,\rm{M}_{\odot}$. |
Based on two epochs of high-accuracy astrometric LIST/WFEPC2 observatious separated by wre. relative proper motions of 829 stars were measured. | Based on two epochs of high-accuracy astrometric /WFPC2 observations separated by yr, relative proper motions of 829 stars were measured. |
A selection of caucdidate cluster members with P4>0.9 results ina clean cluster CMD. | A selection of candidate cluster members with $P_{\rm mem}>0.9$ results in a clean cluster CMD. |
The best-fitting isochrone vields an age of LAINIvy. a distance of kkpe. aud Ay=L6-L7uunae for the PAIS aud intermeciate-mass AIS cluster members. | The best-fitting isochrone yields an age of Myr, a distance of kpc, and $A_{\rm V}$ mag for the PMS and intermediate-mass MS cluster members. |
Stars at the location of the short-lived radiative convective gap. which occurs at AMAL. at the age of NYC. are identified for the first time. | Stars at the location of the short-lived radiative convective gap, which occurs at $_{\odot}$ at the age of NYC, are identified for the first time. |
We find hiuts of a sparse voung low-mass population with an age of [MMyrL which might constitute an earlier generation of star formation im 336023. and likely represeuts the low-nass counterparts to several blue supergiauts iu the vicinity of NYC. | We find hints of a sparse young low-mass population with an age of $\sim 4$ Myr, which might constitute an earlier generation of star formation in 3603, and likely represents the low-mass counterparts to several blue supergiants in the vicinity of NYC. |
For the first time. the internal velocity dispersion of the starburst cluster NYC could be measured. | For the first time, the internal velocity dispersion of the starburst cluster NYC could be measured. |
For stars with masses MAL. M 9ALD... we determine a one-dimensional velocity. dispersion of 111+ 2746s vr|. corresponding to L5+ Oskkin s3 atf a distance of kkpe. | For stars with masses $_\odot<$ $<9\,$ $_\odot$, we determine a one-dimensional velocity dispersion of $141\pm27 \mu$ as $^{-1}$, corresponding to $4.5\pm0.8$ km $^{-1}$ at a distance of kpc. |
From the fact that the velocity dispersion does nof vary with stellar mass in this mass range. we deduce that NYC has not vet reached equipartition of enerev. | From the fact that the velocity dispersion does not vary with stellar mass in this mass range, we deduce that NYC has not yet reached equipartition of energy. |
This is not eutirelv unexpected at the vouug age of the cluster. since its crossing time is estimated to be MM by Nüirrubereer Petr-Gotzeus (2002). | This is not entirely unexpected at the young age of the cluster, since its crossing time is estimated to be Myr by Nürrnberger Petr-Gotzens (2002). |
The same ασ be truce for many extragalactic starburst clusters. where qiass estimates rely on the neasurements of velocity dispersous. | The same might be true for many extragalactic starburst clusters, where mass estimates rely on the measurements of velocity dispersions. |
If these clusters are also not vet in virial equilibriuu. their masses might ve systematically overestimated. | If these clusters are also not yet in virial equilibrium, their masses might be systematically overestimated. |
Thus. NYC provides au iuportaut beuchiuark for our uuderstaudius of the carly dynamical evolution aud the long-term survival of voung. nassive stellar clusters iu the Milkv Way aud in other | Thus, NYC provides an important benchmark for our understanding of the early dynamical evolution and the long-term survival of young, massive stellar clusters in the Milky Way and in other |
that could be inserted. directly into the existing VLA 8.0-8.8 Gllz ervostat with minimal modifications. | that could be inserted directly into the existing VLA 8.0-8.8 GHz cryostat with minimal modifications. |
A planar OMT [16]-|18] may be the only design that will allow this option to be viable. | A planar OMT [16]-[18] may be the only design that will allow this option to be viable. |
This (wpe of design replaces (he coaxial probes with a four probe microstrip circuit Chat requires (wo 180* hybrid couplers to combine the signals from opposing probes. as well as a 90° hvbrid to generate the left and right circular polarizations [rom the two orthogonal linear polarizations. | This type of design replaces the coaxial probes with a four probe microstrip circuit that requires two $180^\circ$ hybrid couplers to combine the signals from opposing probes, as well as a $90^\circ$ hybrid to generate the left and right circular polarizations from the two orthogonal linear polarizations. |
While the benefits of (he small size and weight of a planar OAT approach are obvious. it is vet unclear how its inherent resistive losses will affect its noise performance compared to a more traditional waveguide polarizer. | While the benefits of the small size and weight of a planar OMT approach are obvious, it is yet unclear how its inherent resistive losses will affect its noise performance compared to a more traditional waveguide polarizer. |
Two tvpes of planar OMIT designs will be investigated - one using eokl thin-fihu microstip circuits while the second uses high temperature superconductors. | Two types of planar OMT designs will be investigated - one using gold thin-film microstip circuits while the second uses high temperature superconductors. |
An all-waveguide OMT design is also being developed in the event that the noise contribution in the sienal path prior to Che low-noise amplifiers is excessive. | An all-waveguide OMT design is also being developed in the event that the noise contribution in the signal path prior to the low-noise amplifiers is excessive. |
This device is based on an ultra-thin turnstile design [19]. | This device is based on an ultra-thin turnstile design [19]. |
Hs lateral dimensions would require a new crvostal to be designed to accommodate it. | Its lateral dimensions would require a new cryostat to be designed to accommodate it. |
There are nine different (vpes of low-noise amplifiers (LNAÀs) used on the EVLA. all of which were designed and built by the NRAO’s Central Development Laboratory [20]-[22]. | There are nine different types of low-noise amplifiers (LNAs) used on the EVLA, all of which were designed and built by the NRAO's Central Development Laboratory [20]-[22]. |
The primary active elements in the amplifiers are indium phosphide. heterostructure. field effect (ransistors (InP HETTs) manulactured by TRW (now Northrop-Grumnman Space Technology) under the NASA-lecd Cryogenic ILEMT. Optimization Program managed by the Jet Propulsion Laboratory. and obtained by the NRAO under special agreement. | The primary active elements in the amplifiers are indium phosphide, heterostructure, field effect transistors (InP HFETs) manufactured by TRW (now Northrop-Grumman Space Technology) under the NASA-led Cryogenic HEMT Optimization Program managed by the Jet Propulsion Laboratory, and obtained by the NRAO under special agreement. |
All the TFET devices used in the input stage of the EVLA amplifiers are [rom the so-called Crvo-37 wafer which exhibits record performance when compared to all other mamuacturing runs of similar devices. | All the HFET devices used in the input stage of the EVLA amplifiers are from the so-called Cryo-3" wafer which exhibits record performance when compared to all other manufacturing runs of similar devices. |
In the 4-100 Gllz range. (he cooling of an InP HETT twpically reduces the internal amplifier noise to within 4 to 6 times the quantum limit [23]. | In the 4-100 GHz range, the cooling of an InP HFET typically reduces the internal amplifier noise to within 4 to 6 times the quantum limit [23]. |
The development of coolable LNAs requires detailed knowledge of both the signal and noise models of ILEIZEs al crvogenic temperatures. | The development of coolable LNAs requires detailed knowledge of both the signal and noise models of HFETs at cryogenic temperatures. |
These models have been developed with sufficient accuracy {ο achieve designs with optimal noise bandwidth performance [24].|25]. | These models have been developed with sufficient accuracy to achieve designs with optimal noise bandwidth performance [24],[25]. |
The design of the latest generation of amplifiers used for the EVLA was largely built upon the success of the LNAs created by the NRAO for the Wilkinson Microwave Anisotropy Probe project. | The design of the latest generation of amplifiers used for the EVLA was largely built upon the success of the LNAs created by the NRAO for the Wilkinson Microwave Anisotropy Probe project. |
The use οἱ] chip aud wire” technology allows for the best performing WFET devices to be chosen from the waler to be emploved as (he input stage of the amplilier. (hus ensuring a superior noise ligure while avoiding a repeatability problem that is often observed in devices αἱ crvogenic temperatures [rom different parts of the waler or [rom different wafers. | The use of chip and wire" technology allows for the best performing HFET devices to be chosen from the wafer to be employed as the input stage of the amplifier, thus ensuring a superior noise figure while avoiding a repeatability problem that is often observed in devices at cryogenic temperatures from different parts of the wafer or from different wafers. |
5ingle-ended amplifier designs are used for most of the frequency. bands. with the C. X and Ku-band LNAs having | Single-ended amplifier designs are used for most of the frequency bands, with the C, X and Ku-band LNAs having |
radiation. | radiation. |
The most general assumption for the spectrum of accelerated electrons is the power-law spectrum wilh an exponential eutolf at energy o: The energy cutoff 5; in the spectrum of electrons is determined by the balance between the electron acceleration and cooling times. | The most general assumption for the spectrum of accelerated electrons is the power-law spectrum with an exponential cutoff at energy $\gamma_0$: The energy cutoff $\gamma_0$ in the spectrum of electrons is determined by the balance between the electron acceleration and cooling times. |
The acceleration time of electrons is usually given bv the following general form: llerealter it is assumed (hat the magnetic Ποια is distributed isotropically. e.g. B= /2/3B. | The acceleration time of electrons is usually given by the following general form: Hereafter it is assumed that the magnetic field is distributed isotropically, e.g. $B_{\perp}=\sqrt{2/3}B$ . |
The parameter 1> presents the rate of acceleration which could be the energy-dependent. | The parameter $\eta\geq
1$ presents the rate of acceleration which could be the energy-dependent. |
In different astrophysical environments. 7) remains a rather uncertain model parameter due to highly unknown acceleration mechanism. | In different astrophysical environments, $\eta$ remains a rather uncertain model parameter due to highly unknown acceleration mechanism. |
The svnchrotron cooling {ime is given by The high energy cutoff of svuchrotron emission is given by which depends only on the parameter η. where à;=1/137 is fine-structure constant. | The synchrotron cooling time is given by Then the energy cutoff $\gamma_0$ is determined by the condition $t_{acc}=t_{syn}$: The high energy cutoff of synchrotron emission is given by which depends only on the parameter $\eta$, where $\alpha_f=1/137$ is fine-structure constant. |
Thus. in the regime of acceleration with 7<100. the svnchrotron radiation emitted bv a hiehlv magnetized compact spot can result in an effective production of hard X-vavs that extend to above MeV. bands. | Thus, in the regime of acceleration with $\eta \leq 100$, the synchrotron radiation emitted by a highly magnetized compact spot can result in an effective production of hard X-rays that extend to above MeV bands. |
The non-relativisticor relativistie shock acceleration is most | The non-relativisticor relativistic shock acceleration is most |
the uncertainties; though taken by themselves they do not provide convincing evidence of a metallicity gradient. | the uncertainties; though taken by themselves they do not provide convincing evidence of a metallicity gradient. |
The results summarized in Fig. | The results summarized in Fig. |
2 depend on the assumption that the LMC RR Lyraes follow the same mean period - [Fe/H] relation as the Galactic stars. | 2 depend on the assumption that the LMC RR Lyraes follow the same mean period - [Fe/H] relation as the Galactic stars. |
This assumption cannot be tested directly. | This assumption cannot be tested directly. |
The LMC RR Lyraes with spectroscopic abundances plotted in Fig. | The LMC RR Lyraes with spectroscopic abundances plotted in Fig. |
1 lie with the Galactic points. | 1 lie with the Galactic points. |
However, their small number, their limited range in [Fe/H] and the substantial intrinsic scatter in logPa» at a given metallicity!,, precludes any attempt to derive an independent slope to the relation. | However, their small number, their limited range in [Fe/H] and the substantial intrinsic scatter in $\log P_{ab}$ at a given metallicity, precludes any attempt to derive an independent slope to the relation. |
A referee has suggested that we should consider the case when there is no relation between [Fe/H] and period and that there might in fact be no metallicity gradient for this population. | A referee has suggested that we should consider the case when there is no relation between [Fe/H] and period and that there might in fact be no metallicity gradient for this population. |
Since Fig. | Since Fig. |
2 is simply a scaled version of a mean logP - Rac relation, the gradient in mean logP must then be explained in terms other than [Fe/H]. | 2 is simply a scaled version of a mean $log P$ - $R_{GC}$ relation, the gradient in mean $\log P$ must then be explained in terms other than [Fe/H]. |
A model with a nearly constant mean [Fe/H] for the RR Lyraes, but with an age gradient might be possible. | A model with a nearly constant mean [Fe/H] for the RR Lyraes, but with an age gradient might be possible. |
It will be clear that our favoured model of a metallicity gradient might well also imply an age gradient. | It will be clear that our favoured model of a metallicity gradient might well also imply an age gradient. |
More complex models might be developed. | More complex models might be developed. |
However, they will all involve a systematic radial gradient in RR. Lyrae properties. | However, they will all involve a systematic radial gradient in RR Lyrae properties. |
Battinelli Demers (2005) suggested that in a stellar system the ratio of carbon-rich to oxygen-rich AGB stars was a function of [Fe/H] and a rederivation of this relationship was given by C09 who found: though with considerable scatter as can be seen from fig. | Battinelli Demers (2005) suggested that in a stellar system the ratio of carbon-rich to oxygen-rich AGB stars was a function of [Fe/H] and a rederivation of this relationship was given by C09 who found: though with considerable scatter as can be seen from fig. |
B1 of C09. | B1 of C09. |
Cioni Habing (2003) selected likely LMC AGB stars (i.e. they were brighter than an adopted tip of the RGB) from the DENIS survey in IJK, and divided them into probable C and M stars using colour criteria. | Cioni Habing (2003) selected likely LMC AGB stars (i.e. they were brighter than an adopted tip of the RGB) from the DENIS survey in $IJK$, and divided them into probable C and M stars using colour criteria. |
These data have been used by C09 to study the C/M number ratio and to estimate [Fe/H] using equation 6 as a function of distance from the centre of the LMC. | These data have been used by C09 to study the C/M number ratio and to estimate [Fe/H] using equation 6 as a function of distance from the centre of the LMC. |
In this work, the numbers of likely C and M stars were counted in bins of size 0.04 square degrees over a grid of 100x bins. | In this work, the numbers of likely C and M stars were counted in bins of size 0.04 square degrees over a grid of $100\times 100$ bins. |
Dr Cioni has very kindly placed these counts at our disposal. | Dr Cioni has very kindly placed these counts at our disposal. |
We proceed as in subsection 3 and calculate the value of Rec for each bin and then sum all the likely C and M stars in the bins lying in annuli about the centre. | We proceed as in subsection 3 and calculate the value of $R_{GC}$ for each bin and then sum all the likely C and M stars in the bins lying in annuli about the centre. |
The results are plotted in Fig. | The results are plotted in Fig. |
5. | 5. |
This figure shows a very slight increase in C/M from Rec1 to Rec=4 kpc. | This figure shows a very slight increase in C/M from $R_{GC} =1$ to $R_{GC} = 4$ kpc. |
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