 Good morning, and I'm Chef Tagliente from the same institute of the previous talker, and I will talk about the importance of the neutron data for astrophysics. So I will go, this one is the Darwin set I will treat, so I will give a short introduction about astrophysics. I will talk about the stellar revolution, the SNR processes, and okay, neutron in the lab in this case. I would like to see that slide that my colleague had before, so it is the problem to be the second, I was always second in the life, I was the second child, so I'm sorry. Suffering this situation, okay, we'll talk about the analysis technique, the astrophysical implication, and the conclusion. So astrophysics is a relatively new, so still in the 50s there was a big debate about when and where the nucleus synthesis took place, so still in the 50s there was a gamma was proposing that the nucleus synthesis was related to some event of the origin of the universe, so it was very close to the big bang time, while in the same period oil was proposing that the fabric of the elements of the nucleus synthesis was taking place in the stars. There was no evidence, which was the fabric of the elements, and only in the 52 there was the, let's say the smoking gun, so the spectral line of techniques were observed in some right gynecology. Now the technicium is a special isotope, it doesn't exist on the earth, because it is not stable, and isotopes with longest life time are the 97, 97, 98, which have 2.6 and 4.2 million years, so they're not as much shorter than 8 of the earth, and it's also shorter than 8 of the red gynecology. So one was the evidence that the star, the nucleus synthesis, the nucleus synthesis happens in the stars, so since that moment we knew that the majority of the chemical elements in the universe are produced through the nuclear reaction in the interior of the star. So this paper was the milestone of the nuclear astrophysics, and in 84 the polar was honored by the Nobel Prize, and this one is the Nobel lecture that was given, sorry, in 83, in 83, and this one was the Nobel lecture that you can find on science in 1984. And so now, since then, we know that the elements are made in the star, so how we can know the abundances of the elements? So fundamentally there are two sources of information, they are independent sources of information, and these are the specter of the photosphere of the sun, and also some special kind of meteorites. It is important to notice that the abundances from these two independent sources are in a remarkable argument between them, so better than 10%. So we can talk about solar abundances, and what we notice when we go to plot the abundances against mass number is that these elements, the abundances of these elements are not spread casually in this plot, but they follow some trends, so there are peaks and dips, and these dips and peaks give us some important information about the stellar evolution, or better, the stellar evolution can explain this particular plot. So let's talk about the stellar evolution. The best way to represent the stellar evolution is just to plot the luminosity of the star in function of the surface temperature of the star. Just pay attention that you go from right to left, so from low temperature to high temperature, and yet from low luminosity up to high luminosity. So here in the bottom right you have the cool and faint star, here you have the cool and bright star, hot and bright star, and hot and dim star. Our sun is here. Another important feature when you go to make this plot is that the stars are not spread around, also here there is a trend. So you see that observing a system of star, here are 25,000 stars, you see that most of them are on this line that goes from the bright and hot corner to the cool and dim corner, and also there are some other areas. So there is a trend, so it is important to explain why the star prefer to stay in certain area and not spread around. So in order to understand that, the most important parameter of the star is the mass. So you see here that the bigger is the mass, higher stays in this position, the initial mass of the star, and the life of the star is determined by its initial mass. So we are here. Another example, the mass goes in this way, and so you have the small star, they are quite long life and the life is almost very quiet, and while there is a big star that goes through supernova with big explosion, most of the time the star is spent in this sequence, it is the main sequence that we saw before, where there was the most populated area, and there is this connection to the mass from this way. So heavier is the mass, faster and shorter is its life. And now in the following I will discuss about the difference of the stellar evolution of the different mass, but another important thing is that how is the distribution of the star, always observing a group of star. What we see is that there are much more small star in this range of mass than heavy star. And also here, now we will talk about this mass range. We have the first, this range was just, if you go here, is in this part. So this part are not really star, the mass is not enough to start barring the hydrogen, but they are able to bar the deuterium, so it is why they are different from the biggest planet. We are talking about solar masses with the solar masses. So this one are not real star, but they can bar the deuterium, but they are not considered real star. The real star starts from this point, from when the mass is in this range. These are called red darts, and they are able to start barring the hydrogen, in this case the star create helium, but the mass is not enough to start barring the helium, and the core of the star became an electron that generates gas, so this gas is able to keep, to contrast to the gravity, and so the star is in equilibrium and slowly slowly dies, but this kind of star has a very long life, like about 10 to 12 years. And here is just to sketch, here is our star, this one is a red dart, and this one is at the brown dart that are these kind of objects. The most important star are the one in this range, between 0, 5 and 3 solar masses. This one is a little bit less for what I am going to say, but always important to study their evolution, but I will not spend much time on this, and all the important masses are very heavy, heavy star. So I will give you some description of this mass, more in detail. The star with the initial mass, including between 0, 5 and 3 solar masses, are much more complicated than one week. So before, already they are able to bar hydrogen, and after when the hydrogen is exhausted, they are able to start the helium barring, they cannot go further, they cannot bar carbon and oxygen. But it is important to see, depending on the mass, they have different way to produce energy, to bar the hydrogen. If the mass is lower than 1.5 solar masses, the fusion goes through the PP chain, while if it is bigger than 1.5, it goes through the C and O cycle. This difference make also change completely the structure of the star, and the way that it is transferred out from the star is different. So for mass between 0, 5 and 1.5, the core is radioactive, while the convective envelope is convective, and for mass for star with mass bigger than 1.5 solar masses, it is the opposite. I will talk longer about this kind of star, because it is very important for the nucleosynthesis. Other stars are the nucleosynthesis of the HV element. These other stars in this mass range, they don't contribute much to the HV element, so we don't spend much time, they just sit at the super HV star, are able to go further stage after the helium, so they can start barring carbon. So far, in this mass range between 0, 5 and 11 solar masses, we have three stages of barring, we can have hydrogen, helium and carbon. Now we go to the heaviest one, so the star with initial mass bigger than 11 solar masses, and they can have other three stages of barring. So after the carbon, they can use neon, they can use oxygen and silicon. So this one is the last stage, after this stage you have the core that is rich of iron, and also this one is the electron gas that generates. So there is equilibrium between the pressure of this gas and the gravity. The problem is that the energy is given by this layer of other material that is the shell around the core, so they are still producing material and the ashes, let's say, that go in the core. So the core becomes bigger and bigger. In one point, the pressure of the electron gas is not able anymore to contrast the gravity and the star starts to collapse, starts to contract by the gravity. So this point of density starts to increase and when the density of the core becomes at the same level of the nuclear density, the nuclei and the free nuclear starts feeling the repulsive part of the nuclear force. At this point, the potential nuclear acts like a stick spring, and there is an explosion, this one is a supernova, so we have that material inside the core is rebounded, and while it is going outward, it encounters the material that is falling from the other shell. So this material is exploding, it is essentially neutron and protons, let's say it is bombarding the ashes created on the other material, so there is a kind of shock waves. So the last part of the star that dies as a different evolution can become a neutron star or a black hole. But this star, they last quite short if you compare it to our sun, or we start with the lower mass, but it is funny to see that for this one is the evolution in the Earth's diagram, and you see that for start with the lower mass, the life is almost complicated, so the line is almost complicated, while for this one it is almost straight, although they have a complicated life. So now we know something about the star evolution, and let's say what we can obtain from this plot of the abundances. So we know that the first part, the helium and hydrogen that is the 19% of all the material around all the abundances in the universe is made by Idrosman and helium, and it comes from the Big Bang. So here we have the gap, that is the gap of the helium, lithium and boron, and these other elements come from the different stages of the fusion in the star, so we have the alpha nuclei, and in the end we have the iron peak that comes from the last stage of the fusion, so the silicon bombarding, and after this we cannot go by fusion, because for the binding energy in this point it doesn't increase any more, but it goes down, so you need energy in order to fuse, so you don't have any energy to spend in the star. So at this point we have most of the 99% of the elements, therefore there are still these small parts, this is 1%, cannot be constrained by fusion. So how they are created, this material, this is the helium, sorry, this is the one created by neutron capture, and the fact that the iron neutron is evident from this peak always, a peak that helps us to explain the physics. So these peaks are just around the neutron magic nuclei, and the fact that there are two peaks, it means that there are two different processes of neutron capture, so these processes are called Earth process and Earth process, and so we can see that all the elements, ever the iron, are the results of neutron capture processes. Okay, now I will talk about the neutron capture processes, so we saw that there are two processes, and essentially there is the Earth process, we would have a lifetime of 10 to the 4 years, and it requires neutron density about 10 to the 8th neutron per cubic centimeter. After we have the Earth process, it happens during the supernova explosion, and you see the lifetime is much shorter, but the density of the neutron is much higher, and between these processes there is a competitor, there is the beta decay, and sometimes there is a few hours, and sometimes it can be in competition, especially with the Earth process, how the Earth process proceeds. So we have these cisnuclei, the ones that are created by fusion in the start, and they can be bombarded or decayed by neutrons, in some way so they capture these neutrons, became heavier and heavier in one point, you have the beta decay, and so you create a new species, and it goes, the path goes in this way, creating all the elements that go from the iron to the light, and you create this element in this beta stability valley. About the Earth process, it goes in a different way, it acts in the supernova explosion, the supernova, you have lots of neutrons, and the density is quite high, and so really, so in this moment when the start is exploding, we have 1.5 giga helping, so all the elements become stronger, they are rich over neutrons, and after they are unstable and they start to decay, they have a chain of decay, still they reach the beta valley stability. So we have that most of these isotopes created by our process are the one close to the neutron decline. The process is stopped. It's important also to note that, okay, behind the iron peak, the elements are created by S and our process, so we have 50% of the contributors to process almost equally, but it's important to note that there are some isotopes that are only S, and they are very important in order to model the star evolution, especially this phase of the nucleosynthesis. In this case, you see that these two, osmium, are S only, because they are shielded by the rhenium 187. So what information we can have from this kind of isotope, especially by the S, I will talk more about the S process than our process because the one that can be described in the laboratory because first of all, you have the daughter, nuclei, it's very related to the parents, it's not the same for the R process, and also you easily, so you would see in the future that the flux of neutron is possible to have in the lab and in the cross-section as a big influence on the full path of the S process. So in the case of the Vier, we have the plot of the cross-section for the abundances for only S isotopes, and so you here see that there is a kind of trend and the value of this sigma by N can change over a factor of 100, but the nice thing is that you see in these two areas between the neutron magic nuclei, we have almost a constant area. In this case, we'll talk about the local equilibrium where we have that the cross-section for the abundances is a constant. It is almost logical because if you think that if you have a small cross-section, of course this nuclei doesn't accept to capture other neutron, so it accumulates. If the cross-section is bigger, it's easy to capture and so the abundances is smaller, and it is logical. So the idea from here, you can see that of course the nuclear synthesis of the element heavier than iron cannot come only from one explosion, so it cannot come only from one star just because you have this kind of exponential trend. So it means that to obtain this material, these elements, they have to go through different generation of the star. So the idea was to describe all these points by an exponential using a few parameters, just the fraction of nuclear seeds and the time of the explosion. So this one was a big success in the 90s. This is the fraction of the seismic nuclei and the tau's neutron exposure. So in the 90s, it was a big success by Canberra. He was able to fit this data using the cross-section and the abundances available in that time. So what he got was that there are essentially two areas, this one is the blue one is the main component and it can be described with two parameters, set of parameters, and there is this part in black one that is called the component that is described from this. So what does it mean? That to obtain the element that goes from 90 to 205, you need 6% of nuclear seeds of iron and they have to capture an average of 10 neutrons per time for any exposure. And in the week you have a larger fraction with less time of exposure. You have the seismic nuclei. So another review about this paper is always by Kepler who was giving. This is very nice to read. It's important if you're interested to read. But this simple model doesn't give many information. Okay, now we know that there are different stellar sites where this nuclear synthesis takes place, but we don't have any information about the flux of the neutron, who is producing the neutron, and the time. So along we need to have these elements. And so almost at the same time, in the same period, there was a model, it's called thermal pulse model, where it was explaining in more detail how this nuclear synthesis takes place in some stars. So from now, what we know from this model is that the stellar process sites are the AGB star, the one with the low AGB star with the initial mass between 1.5 and 3 solar masses. This one is for the main component. And instead for the weak component, we have the factory, these heavy star between 50 and 25 solar masses. And in particular during the helium burning in the core phase and in the shell carbon burning. The most attractive between these two is, okay, this one, and I will describe in detail, and here you have two different sources, and I will explain what's happening in this kind of star. So this one is a path in the Erzplan-Grassel diagram. And so the star that we are interested are in this point, usually, this is the mass region that we are interested in for a two solar masses star. And here is when it is the main sequence. When the hydrangea in the core is exhausted, the star is in this position. At this point, the hydrangea is not completely exhausted, but there is not anymore in the core, so the core is helium rich. And so the temperature is not enough to start the helium burning, and the star starts to contract. But still, there is part of the energy that is given from the hydrogen that is in the shell outside the core. But, okay, so we see that the star starts to cool down a bit, but becomes brighter because of course the hydrogen shell is burning. At one point, when the temperature is enough, so because, okay, there is still the contraction for the gravity, and the temperature in the center is increasing. So at one point, the temperature is high enough to start the helium burning. And so the star is in this position at the moment, and so when the helium is burning, so it starts to, again, become hotter, the star, and it moves in this position. So it is a bit less bright, but it is hotter. At this point, the helium is exhausted, and the star has an oxygen-carbon core and two shells outside this core. One is the first one, near the core is helium, and one little bit further is the hydrogen. And the mass of the star is not enough to start the carbon burning, and so you have, again, the contraction. Well, if you have this contraction, the temperature increases a bit, and so it becomes enough to start, again, the hydrogen burning in the shell. The energy that is produced, it is not enough to stop the contraction, but in this moment of the star, you have that, okay, it starts to produce neutron, but okay, you have still the contraction, the temperature increases, and at this point, at one point, you have the enough, the temperature is high enough to start the helium burning. When the helium starts, on the shell always, when the helium starts burning, there is an explosion. The hydrogen shell is brought far away from the core, so it cools down completely, the hydrogen burning, and you have the helium. The helium, okay, it is not just burning, but okay, it doesn't stop the contraction anyway, and after a while, it doesn't have any more, it's not any more enough to contrast the gravity, and you have the hydrogen comes back, the helium switches off, and you have, again, the helium. So you have this kind of a column phase of the star that has contraction and expansion. In this particular part of the life of the star, the star loses the 50% of its mass, and so slowly, slowly, here it is, okay, slowly, slowly, it goes to finish like white, dark. It's important also to see that they are called AGB because you see this one, they tend asymptotically to the giant branch. So let's see what happened. This part is the most important part for the nucleosynthesis of the heavy element. So the star looks like this one. So this one is all the material that is spread from the star, and the real star is this one in the center. So we have the convective envelope, and the star, we say we have this core of carbon oxygen. Here we have the helium barring shell, and here is the hydrogen barring shell. So this one is more, just to see a little bit better, but what happened is very important. So during the hydrogen phase, we have this source, there are lots of protons available, so this source of a neutron is activated. So the flux of this neutron density goes like this. This phase lasts almost 100,000 years. After when the temperature is enough, you have this explosion, and you have the helium barring time, these other sources are activated. And this one you see has much higher neutron densities, but the time is much shorter. It's around 200 years. And these pulse, they are called thermal pulses, they can last between 10 and 15 of these pulses, depending on the mass of the star. So these are the two sources of the neutron. You know, the important thing is during this explosion, you have the part of the material that is in the core, and here it is brought up, so it can refill the carbon here, and part of the heavy elements there, synthesized in this phase, are brought by the solar wind far away from the core. So that's what the outlet you see around the star. Okay, now it's finished. Let's say the spherical part is going out in the lab, and here in this part, I will repeat a bit of just a few things, a few slides of Nicola. So okay, you can have a neutron in the lab, like already we saw with the thermal reactor, or with some energetic neutron, based on this reaction. Nicola was already detailed on this description, or by time, in the time of facility, where you have a wide neutron energy spectrum and high energy resolution. Okay, I think I can go very fast, so there is a different way to produce neutron. One is with the photo-production, and the facility that uses this technique are Aurela, Gelina, and RBI. Okay, Aurela? I think it is not so effective. And while the spallation, you have that neutron are ejected from heavy target due to charged particles impinging on this target facility that uses this technique are end-of-lamps and the eases. The time you already see this one, so this one, so you can calculate the energy from the time of flight. And one of the things that is important for the resolution is the flight path. So longer is the flight path, better is the resolution, but longer is the flight path. The smaller is the neutron flux, so you have to compromise between these two needs. And, okay, this one is in top. You already saw this one before, so I don't need to show you. And this one is Gelina. You see all this one, the flight path of Gelina. You have a better description later by the people from Gelina. And the most important thing is that they have nine flight paths and they can make different kind of measurement. You see here the list of measurement that they can make. About flux, we use already for the integrated flux, end-of is lower than the other facility. Also, you have to think that these flux is at 105 meters. While we go on a instantaneous neutron flux, it's comparable to length, but length is 20 meters and end-of is 185 meters. And this one is very important when you are measuring radioactive samples that are also important for astrophysics. Another important thing that was told in the previous talk was about the resolution. And usually when you have a neutron, it can escape directly or you can have a path inside the spallation target. So what you have is you have the geometrical length plus the path inside of the sample. So what you can happen is, you have to take account of the length because you can have a neutron that arrives at the same time, but they have originally different energy. So in this case, talk about resolution and in order to take account of the resolution, you have this way, it depends on the energy in this way. So you have to take account of this time-off. What is the effect? You see here is the comparison between gelina and end-of. You see that the effect of the resolution in the gelina is much smaller than one in end-of. You see that how it depends on the distance. So in the end-of, there is a strong influence of the resolution. What is the effect of the resolution? So this one is what you should see. Resonance, this one is the product due to the Doppler effect and when you consider also the resolution, this one is what you really have. So all this effect has to be considered in the data analysis when you make the analysis of the resonance. What do we detect? Usually for astrophysical needs, you have two different kinds of measurements you can make. One is the capture. So you have the neutron arriving of the sample. The capture, they emit a gamma, so you detect the gamma and you can have also transmission measurement where here you measure all the neutron that doesn't, the neutrons that don't undergo to the capture. What about the neutron capture? What happens is that a neutron is captured, so you have this compound nucleon. So after the compound nucleon stays in some excited level, it decays emitting gammas from the detection of the gamma, you have the spectra of the resonances. So what we measure, we measure the cross-section of the compound, this one is the formula over here, the meaning of all these samples, the spin factor is G and the gamma is the total width and the neutron width. And where the total width usually is given by all the exit channels, so by the elastic channel, the capture channel and the fission. Most of the sample of the target for astrophysical interest for most of the sample is the gamma f can be neglected. So if you see a resonance, this one is the way the meaning of the gamma, this one is the highest proportion of the gamma gamma gamma and the total gamma, and one is the energy of the resonance and the total relative area is connected to these two parameters, the gamma and the gamma gamma. Another important thing is when you're detecting, so, you have different resonances, so there is a definition, that we have the resonance region where your gamma resonance is smaller than the distance between two resonances and in the resonance region, you have the resolution of the resonant region when you have the detection of resolution is smaller than this distance and you have the unresolved resonance region when the detection resolution is bigger. So, we will talk about the different techniques, how you can detect the gamma, one is from activation, first of all you bombard the samples that you want to study and after you go to detect all the gamma that are produced by this sample and it cannot be done for all the samples because it can be done only for samples that have the A plus one isotope radioactive otherwise it is stable, it cannot be done and there is no time of flight after another way to study it in your capture is by the level population it is visible with the germanium detector and the other one of the techniques is the total energy detection where you detect only one gamma from the cascade so your detector has such a small efficiency that it is able to detect just one gamma per time in this case you need to do it in function that we will explain later and the detector has to be insensitive to the neutron and typical detector for this kind of detection are the C66 that we are using in Galena and Tov another way is the one that when you detect all the cascade so in this case you need a 4 pi detector with high efficiency very high efficiency and it allows to have discrimination between capture and fission typical detector for this kind of spectroscopy are the barium fluoride and they are used in Tov let's see now what is the yield most of you know so the yield is just a fraction of instant neutrons that undergoes to an n-gamma reaction so here is the number of neutrons that undergoes this reaction this one is the total flux and this one is the efficiency and the relation between yield and cross section is reported here important things is the efficiency is one that like you see it depends from the neutron energy from the energy of the neutron and the energy deposited in the detector so we saw before that one way to detect the neutron was the total energy detection where you were detecting only one gamma and so here you have to pay attention because of course you can detect this gamma or you can detect this gamma they have different energies so the efficiency will be different for the same neutron so what you have to do is just to find a way in order to have this efficiency independent of the deposit energy so what we do that with the pulse wave technique so just simply what you do you just modify the software of your efficiency in order to make your efficiency independent from the energy of the single gamma so the independent from the cascade you do that making simulations you just take in mind that you know the energy of the neutron by the time of flight so you know the energy of the neutron when you detect the gamma you use the right way in order to have the efficiency independent from the cascade and the accuracy of this technique is better than 2% another important thing that we saw before is the fact that these detectors have to be insensitive to the neutrons but what can happen even if you pay attention if you like C66 are almost they don't detect the neutrons but usually what you have is this one so you have a neutron that arrives and the gamma is produced and one of the gamma is detected from your C66 but what can happen is, so you have this spectra so what can happen is that the neutron is elapsed scattered and after it's captured by the material surrounding the detector and the gamma that is produced by this capture is confused with one capture event and so in this case you have this position of these two the elastic with capture cross section so you have to pay attention on the material that is surrounding or the material which your detector are made in the case like many cases we have the detector are made in carbon fiber and most of Natalia that came where the capture can occur is avoided so it is not used another important thing that has to be experimentally has to pay attention has to pay attention is the thickness of the sample because as you can see you can have a sample single scattering so here you have them the right one is the resonance within the single scattering but you can have the sample is too big you can have multiple scattering and you can see that the resolution of the resonance get worse in this case and so it's just now about the analysis so before that just because of the instrumental limitation we have that our resolution get worse of the Doppler effect and the resolution of them get worse but the only important thing is that the area of all these resonance is the same so the area below the resonance is independent of the experimental resolution and so what is important to study more than the resonance is the area and it's already before the relation this one is the relation between the total area and the cross section this one is the relation between the gamma capture and the cross section and the area for the capture so in the case of the thing sample we have that for the transmission we can express the area with this formula and you pay attention that it's different if you use thin and thick sample in the case of thin sample from this area, from that area analysis you can detect, you can have, sorry, directly the gamma N in the case of thick you have this relation between the area and the gammas in the case of the capture in this case you have that you have to use always thin sample and the relation between the area and the parameters gamma N and gamma gamma is this one so it is important to note that for samples where the gamma N is much bigger than the gamma gamma from the area you can directly calculate the gamma gamma but in case it is not possible if you measure only with the capture what you can have is just this relation where you can calculate this factor it's called kernel and it can be used for astrophysical so ok, this one are all the techniques that we can use for a calculator cross-section and how we use the cross-section for the astrophysics so what is important for important things is normalization so we saw before that this efficiency depends on the neutron energy and on the energy that is detected but also it depends from other other parameters that are difficult to keep under control and these parameters are like the liquid of the circulator that is changed with temperature or the distance between the circulators and the detector that can change just when somebody changed the sample so to keep under control to take into account these other factors one is to normalize the measurement so you make a measurement of some standard sample in this case it's the gold but you can use also the iron so there are some samples that are very well known you use the saturated resonance and what you do you calculate the yield for this parameter because you know very well this parameter you make a fit with your theoretical data and so you determine these two parameters and so this one is the normalization factor that you have to use in order to take into account all these parameters ok now we are here so we saw how to calculate the cross-section which we know that in our case the area of the resonance is more important so to normalize but what we need for astrophysics we know that the neutron that are producing the star are quickly thermalized to the temperature of the star and then so what you have to use is the martian Boltzmann distribution and so you have to follow the cross-section that you have for your sample that you have studied with the masculine average cross-section of the temperature that you have studied so we saw before from the AGB star that there are different temperatures so during the carbon-13 neutron source the temperature is around 8 kV for the neon neutron source the temperature is around 23 kV while for the as process in the AGB star the temperature is around 90 kV so you see that you need to calculate your maximum average cross-section for different macwellium so what usually you do you calculate the maximum average cross-section in this way using all the cross-section the range of folding the cross-section with the maximum average with the maximum distribution ok so how you can calculate the maximum average cross-section depends on which facility you work so when you use the activation technique you have this more energetic almost more energetic neutron spark they are already this maximum shape spectra and so the data can be directly the data that you obtain from this facility are already averaged for the maximum distribution facilities like this are ok usually in the vandegrafs there are facilities like SARA, FRANCE, I think SARA is already producing data and so this measurement is in two steps you have the activation and the second you measure the gammas the pro is that it is a selective method so there is a good ratio in signal and background and you have a larger flux because the sample is just attached to the the neutron source the problem is that you can use for all the samples and you also think that you can produce the maximum average cross-section only for certain energy it can be 25 or 30 kV so it is not easy to extrapolate this data to lower or higher energy especially for samples that presents lots of resonance in the region that you want to study so it is not so straightforward the other one is with the neutron time of light we saw that quite a lot so here you can determine the cross-section for almost all energy so you can calculate maximum average for all energy the problem is that essentially you need almost pure sample while in the previous when you use activation you can use just natural samples and also sometimes they have to be quite big also here we have the problem but most of the time it is just the samples the biggest problem and so when you have a negative isotope you need a very high flux facility ok now I will show you the needs for the astrophysics which cross-section do we need so here is the almost actual situation about accuracy and the mass number so you see that the accuracy of the cross-section is quite high so it is quite low for most of the samples and what is required in astrophysics is an accuracy better than 5% with the control of the systematic so you can see that there is lots of work to do between these elements there are some that are more important like the branching point where the beta decay is in competition with the S process after you have the much nuclei and points like the lead the seed isotopes are very important to know in order to modeling the stellar evolution and there is some isotopes of special interest like the cross-mochronometer or neutron poison isotope about the branching point along the S process there are 21, this one is a nice paper that talks about the neurological physics in the start and here there are the list of the branching points, there are 21 on the S process path and before 2001 no one of them was measured so recently since 2001 two of them were measured the nickel 63 and the samarium 150 three of them were measured this year but still there are not the results the complication of these samples is not the fact that they are radioactive ok, you see here there is the life time so for example for the nickel it's 100 years and for the samarium it's done 900 here the promensium is 2 years and this one is 1.9 so it's not only the radioactive but how to obtain these samples is not easy to find them and also to use them in the facility so this one the complication comes just from the samples how to obtain them so I will describe a bit the samarium, the samarium is very important because it has half life, it's 90 years on the laboratory but it's changed with the temperature so in the S process site it can go down till 3 years and so depending on this temperature you can have different flux of course you can have to go on this path or can go on this other way so it can fit european 151 european 153 in different way so it will follow different path so the branching rate depends on the thermal condition of the star and it can be used to measure the temperature of the star of course once the cross section is well known this experiment was done at end of 2001 and for the first time it was possible to see many, many resonance and here is the background here is the signal so you can see with this kind of facility what kind of measurement you can make with the radioactive samples and what was for the first time we were able to have the experimental maximum level cross section you can see that here in the yellow you have the models experimental one so there was big difference between the models and the real data another important isotope that was measured another branching point was the nickel 63 this one is one of the seeds nuclei just in this region for the masses smaller than 90 and so in this here this nickel 63 is the first branching point of the S process because you of course you can have that depending on the flux you can have that one path or the other and so you have that during the helium core burning the S process proceeds in this way and during the carbon shock burning the process goes in this way another important things is that for this region for the mass smaller than 90 the relation sigma and constant is not valid so the local approximation is not valid so it is important to know very well the cross section because the accuracy of one cross section influence the knowledge of the other the other cross section well also here just to give an idea how difficult it is to have the samples is the sample of the nickel 63 was obtained from the nickel 62 sample it was irradiated in a thermal reactor for a few years and in the end we obtained only 30% of the nickel 63 you have to take into account the decays and so there was also a contamination from by coolio 63 and in this case it was possible to make the payment of separation to obtain almost a copper free sample but for many same case it is not so easy so in the end we obtained these results and also this one is the maximum experimental maximum average cross section that also in this case is different from the one that was in the models before so it was the first high resolution measurement need for this sample another case important for the astrophysics is the zirconium the zirconium is a surrogate of the diamond but it is very important in astrophysics and it belongs to the first Neutromagic the bottleneck for the S process like we say before and so we have just it is important for the HGP star so see here in the first phase during the carbon phase you have that S process goes in this way so the 96 is completely steep from the production of the Neutron capture and why when you have the higher Neutron flux you have the production of 96 is important to note that in the second phase of the pathos so when you have this one is the first phase hydrogen helium so after when you have again hydrogen the zirconium 96 is not reached by the S process but it is destroyed by the S process because it is there so it captures the Neutron and after it is created again so the quantity cross section and the abundance of the zirconium 96 is very important to constrain the stellar model and also it is important because after 1.5 mega here you have that the zirconium 93 big case of the neobium 93 that is the only stable isotope of the neobium so the importance is also ok what was the situation just in the past there were many measurements during the 80 sorry there were many measurements of this isotope and what was noticed was that the zirconium 90 here we have the solar abundances and here you have the ratio between what you calculate and you calculate with the cross section so the S abundances over the solar abundances so it means that the S process contributed to the 80% of the zirconium 90 while the zirconium 90 should be S only isotope and here the 96 should be almost only air isotope because it partially is escaped from the S process so it is produced only during the neon phase but after it is destroyed by the hydrogen phase carbon phase that is much longer you see that here the abundances due to the S process calculated with the old cross section was higher than the one from the zirconium 90 which should be much smaller the importance of the zirconium is also in the meteorites when you study them and you see that the zirconium 96 is the one that has bigger fluctuation because it gives us an indication of how many pulses there were the thermal pulses there were in the star and so on it is very important for the S process and the data that were obtained in this measurement where this one is with the old cross section that were used to make the calculation we saw before this part of 95 is missed because this is a branching point with a very short life so it is not easy to measure so the only way that you have to guess about this the estimation about this cross section is using some model or using some kind of calculation so what we did in the analysis was just to consider all the maximum for the even isotopes how they were varying in function of the neutron number this one is for the you see so this behavior is just because it is a neutron magic so here the cross section is very low because the neutron magic actually starts to increase and decrease while if you go a little bit far away from the region of the neutron magic you have that the trend is more linear so we took the odds also and what we try to do is to you make a fit of these two lines and the only difference between the and even isotope is the pairing effect so this way you can have an estimation of the cross section of the 95 so what you was calculating this way was like 21 millibar and if you compare with models you see this was not too far from this model most so with this new value we were to calculate again the ratio between as abundances and so abundances that you see that this the guarantee is much better but finally you can see that the guarantee is almost free so the contribution from this process is as it should be so this one is the results are written in this paper another okay I'm still okay now at the time okay so I will finish it so it is another important so we said that there are some isotopes that can give you some information about the age of the universe so the Osmo is one of these isotopes and so with that we can know how old is the the universe and just okay the data are here in the so this was a measurement we published here and just you see how we can measure the age of the universe one way is the cosmological way by the time definition the other one is the astronomical way based on the observation of global class another way is the nuclear way that it is based on the abundances and decay properties of the long-life species there are different isotopes that can be used for this but notice that the only one that is in the S path is this where the isotopes are smaller than the left because the S process goes from the iron to the iron are this Osmo and the other one can be used they are in the are created by a process so the idea is that okay here the idea is that we have this one RS only isotopes and the renewal is produced by by our process and it decays after 41 giga here so if you know very well the cross section of 186 and 187 you can the contribution to 187 due to the decay and from this contribution you can calculate the age of the universe so the idea is that you are in the local approximation so we are here in this region where n is constant and you have so that sigma 96 times the abundances of 86 as to be equal to the cross section of 87 for the abundances of 87 so in this way you can have the contribution due to the decay of the can be calculated in this way so what you need is just to calculate these abundances astronomical observation and so you need to calculate this very well this section of 86 87 this measurement was done also at the end of this one is the cross section for the maximum average cross section for the 86 and you see this difference was smaller than the previous measurement and this one is for the 87 and in this case you have that the maximum average cross section is a bit bigger than previous measurement in this way they were able to calculate the age of the universe and they were given this way the most important thing is the deduction of the error that now we are at the cross section only 0.4% to 0.4 Giga to the total error and so you see that there is still a difference between all these techniques so this way to calculate the nuclear the age of the universe should match in the same range and ok there can be some problem for the other but of course we know the problem for our cross section is that a part of the the fortune condition for the osmium we have that there are many problems like the beta decay of life of the region is strongly dependent on the temperature of the the stellar electron the capital cross section of the osmium 87 is influenced by the the population of the low line excited level and also ok the branching of Tungsten and Renium can be also some contribute to the to the complication of this measurement and also the chemical evolution of the galaxy influenced the history of the nuclear synthesis so in this way also I can conclude now so I think I show you how much is important to use the cross section for astrophysics and the last 5 years there was a big progress there was a first part of just after that the 60 there was a a big number of measurement that were made for astrophysical purpose after the 80 there are measurements now recently also to the end of facility that this interest is renewed and so there were many interest and many cross section were calculated also interesting cross section like the one I showed a few of them I showed here but still we need further progress especially in the branching point where they are very difficult to measure so it is one of the challenge for the future ok thanks